{
    "0201/astro-ph0201411_arXiv.txt": {
        "abstract": "Short period double degenerates (DDs) are close white~dwarf\\,--\\,white~dwarf binary stars which are the result of the evolution of interacting binary stars. We present the first definitive measurements of the mass ratio for two DDs, WD\\,0136+768 and WD\\,1204+450, and an improved measurement of the mass ratio for WD\\,0957$-$666. We compare the properties of the 6 known DDs with measured mass ratios to the predictions of various theoretical models. We confirm the result that standard models for the formation of DDs do not predict sufficient DDs with mass ratios near 1. We also show that the observed difference in cooling ages between white dwarfs in DDs is a useful constraint on the initial mass ratio of the binary. A more careful analysis of the properties of the white dwarf pair WD\\,1704+481.2 leads us to conclude that the brighter white dwarf is older than its fainter companion. This is the opposite of the usual case for DDs and is caused by the more massive white dwarf being smaller and cooling faster. The mass ratio in the sense (mass of younger star)/(mass of older star) is then 1.43$\\pm$0.06 rather than the value 0.70$\\pm$0.03 given previously. ",
        "introduction": "Short period double degenerates (DDs) are binary stars in which both stars are white dwarfs. The orbital periods of DDs are hours or days (Maxted \\& Marsh 1999) so the separation of the stars is only a few solar radii.  White dwarfs are the remnants of giant stars which have radii of hundreds of solar radii, much larger than the current size of the binary, so there has clearly been dramatic shrinkage of the binary orbit during the evolution of the binary star. There are several models for the formation of DDs which have been used to predict the properties of this group of binary stars. These usually assume that the most recent episode of orbital shrinkage is due to a common envelope (CE) phase  in which a red giant star comes into contact with its Roche lobe and begins to transfer mass to its companion star. This mass transfer is highly unstable, so a ``common envelope'' forms around the companion and the core of the red giant. The drag on the companion orbiting inside the common envelope  leads to extensive mass loss and dramatic shrinkage of the orbit (Iben \\& Livio 1993). If this CE phase happens while the star is on the first giant branch, the degenerate helium core of the red giant will be exposed and will appear as a white dwarf of unusually low mass, i.e., about 0.4\\Msolar, c.f., 0.55\\,--\\,0.6\\Msolar for a typical white dwarf (Bragaglia et~al. 1995). The mass of a white dwarf can be measured directly from its spectrum by comparing the surface gravity, $\\log g$, and the effective temperature, T$_{\\rm eff}$, to cooling models of white dwarfs. Low mass white dwarfs identified this way are particularly fruitful source of DDs (Marsh, Dhillon \\& Duck 1995).  The most notable difference between models for the formation of DDs is the way they treat the formation of the first white dwarf. For example, Iben, Tutukov \\& Yungelson (1997) have predicted the distributions of masses, periods and mass ratios for DDs using a numerical model of the population of close binaries which evolve through either two CE phases or an Algol-like phase (stable mass transfer on a thermal timescale) followed by a CE phase. Han (1998) has produced a similar model but has also explored how the various parameters of the model affect the distribution of periods, mass or mass ratio which are predicted. Han's model also includes enhanced mass loss in a star as it approaches its Roche limit, an effect which is not incorporated in the model of Iben, Tutukov \\& Yungelson. It has now been fairly well established that these models do not succesfully predict the distribution of mass ratios for DDs. This problem has been explored on a case-by-case basis by Nelemans et~al. (2000). The size of a red giant is related directly to its core mass so they were able to show that for three DDs with measured mass ratios, the standard prescriptions for the first mass transfer phase predict orbits which are too small. They used a parametric approach to describe the first mass transfer phase based on a consideration of the angular momentum balance during this phase, rather than the more usual energy balance arguments used by Iben et~al. and Han et~al. This parametric approach was incorporated into a model for the population of DDs by Nelemans et~al. (2001). They were able to predict a mass ratio distribution which was more consistent with those observed provided that they assumed that very low mass white dwarfs ($\\la 0.3$\\Msolar) cool more rapidly than recent models predict. This assumption is required to avoid the prediction that most binary white dwarfs should have very low masses - a problem common to many of these models.  Measuring  mass ratios for DDs is a powerful way to test models of how binary stars interact. Strong obervational tests of these models are desirable because many interesting astrophysical phenomena are the result of interacting binary stars, e.g., black hole binaries, AM CVn binaries, Type~Ia supernovae, cataclysmic variables and novae. The properties of these objects have a direct bearing on areas of astronomy other than the study of binary stars themselves, e.g., the evolution of the properties Type~Ia supernovae over the history of the Universe is a matter of direct concern when they are used as standard candles to measure cosmological parameters (Umeda et~al. 1999).  Several preliminary estimates of the mass ratios and others parameters for WD\\,0136+768, WD\\,0957$-$666 and  WD\\,1204+450 have been published elsewhere (e.g., Moran, Marsh \\& Maxted 1999).  The values given here should be used in preference to those earlier estimates although the conclusions discussed above, many of which are based on those preliminary estimates, are not affected in general by the small changes to the values of the mass ratio given here. ",
        "conclusions": "We have determined the mass ratios and orbital periods of the three double degenerate stars WD\\,0136+768, WD\\,1204+450 and WD\\,0957$-$666 by measuring the  radial velocities of both stars from the sharp cores of the H$\\alpha$ line. We have compared the measured mass ratios and orbital periods for these DDs and three others to two models for the formation of DDs, including selection effects, from Neleman's et~al. (2001). We confirm the result that standard models for the formation of DDs do not predict suffucient DDs with mass ratios near 1. We have also shown that the observed difference in cooling ages between white dwarfs in DDs is a useful constraint on the initial mass ratio of the binary."
    },
    "0201/astro-ph0201282_arXiv.txt": {
        "abstract": "The gamma-ray burst \\grb\\ had the highest $\\gamma$-ray peak flux of any  event localized by {\\em BeppoSAX} as yet but it did not have a detected optical afterglow, despite prompt and deep searches down to $R_{lim}\\approx 23.5$. It is therefore one of the events recently classified as dark GRBs, whose origin is still unclear.  {\\em Chandra} observations allowed us to localize the X-ray afterglow of \\grb\\ to within $\\approx 1$\\arcsec\\ and a radio transient was detected with the VLA. The precise X-ray and radio positions allowed us to identify the likely host galaxy of this burst, and to measure its redshift, $z=0.846$. The probability that this galaxy is a field object is $\\approx 1.6\\times 10^{-2}$. The X-ray spectrum of the afterglow shows significant absorption in excess of the Galactic one corresponding, at the redshift of the galaxy, to $N_H=(5\\pm1)\\times 10^{21}$ cm$^{-2}$.  The amount of dust needed to absorb the optical flux of this object is consistent with the above HI column density, given a dust-to-gas ratio similar to that of our Galaxy.  We do not find evidence for a partially ionized absorber expected if the absorption takes place in a Giant Molecular Cloud. We therefore conclude that either the gas is local to the GRB, but is condensed in small-scale high-density ($n\\gtrsim10^9$ cm$^{-3}$) clouds, or that the GRB is located in a dusty, gas-rich region of the galaxy.  Finally, we examine the hypothesis that \\grb\\ lies at $z\\gtrsim5$ (and therefore that the optical flux is extinguished by Ly$_\\alpha$ forest clouds), but we conclude that the X-ray absorbing medium would have to be substantially thicker from that observed in GRBs with optical afterglows. ",
        "introduction": "\\label{sec:intro}  It is observationally well-established that about half of accurately localized  gamma-ray bursts (GRBs) do not produce a detectable optical afterglow \\citep{fkw+00,fjg+01}, while most of them ($\\approx90\\%$) have an X-ray afterglow \\citep{piro01}. Statistical studies have shown that the optical searches of these events, known variously as ``dark GRBs'', ``failed optical afterglows'' (FOA), or ``gamma-ray bursts hiding an optical source-transient'' (GHOST), have been carried out to magnitude limits fainter on average than the  known sample of optical afterglows (\\citet{lcg01,ry01}, but see also \\citet{fjg+01}). Some of these GRBs could be intrinsically faint events, but this fraction cannot be very high, because the majority of dark GRBs shows the presence of an X-ray afterglow similar to that observed in GRBs with optical afterglows \\citep{piro01,lcg01}. Thus dark bursts could constitute a distinct class of events and not only be the result of an inadequate optical search, but it is unclear whether this observational property derives from a  single origin or it is a combination of different causes.  If the progenitors of long-duration GRBs are massive stars \\citep{pac98b}, as current evidence suggests ({\\em e.g.,} \\citet{bkd+99,pgg+00}), extinction of optical flux by dusty star-forming regions is likely to occur for a substantial fraction of events ({\\it the obscuration scenario}). Another possibility is that dark GRBs are located at redshift $z\\gsim$5, with the optical flux being absorbed by the intervening Ly$\\alpha$ forest clouds ({\\it the high-redshift scenario}).  Dark bursts which can be localized to arcsecond accuracy, through a detection of either   their X-ray or radio afterglow, are of particular interest.  The first and best-studied example was GRB\\,970828 for which prompt, deep searches down to R$\\sim$24.5 failed to detect an optical afterglow \\citep{odk+97, ggv+98d} despite it was localized within a region of only 10\\arcsec\\ radius by the ROSAT satellite \\citep{gse+97}. \\citet{dfk+01} recently showed how the detection of a short-lived radio transient for GRB\\, 970828 allowed them to identify the probable host galaxy and to infer its properties (redshift, luminosity and morphology). In addition, they used estimates of the column density of absorbing gas from X-ray data, and lower limits on the rest frame extinction (A$_V>3.8$) to quantify the amount of obscuration towards the GRB.   Given the extreme luminosity of GRBs and their probable association with massive stars, it is expected that some fraction of events will be located beyond $z>5$ \\citep{lr00}. These would be probably classified as dark bursts because the UV light, which is strongly attenuated by absorption in the Ly$_\\alpha$ forest, is redshifted into the optical band. Fruchter (1999)\\nocite{fru99} first suggested such an explanation for the extreme red color of the optical/NIR emission for GRB\\,980329, although an alternative explanation based on   H$_2$ absorption in the GRB environment would imply a somewhat lower redshift \\citep{draine00}. In a recent paper \\citet{jah+02} derive a photometric redshift $z\\approx3.5$. We note that the three redshifts determined or suggested so far for dark GRBs ($z=0.96$, GRB970828, \\citet{dfk+01};$z=1.3$, GRB990506, \\citet{tbf+00, bkd02}; $z\\approx0.47$, GRB000214, \\citet{apv+00}) are in the range of those measured for most bright optical afterglows, but whether this applies to the majority of these events is still to be assessed. Particularly interesting in this respect is the case of the  so-called X-ray flashes or  X-ray rich GRBs discovered by {\\em BeppoSAX}\\citep{hzkw01}. In most of these events no optical counterpart has been found. The only tentative association claimed as yet is for the event of Oct.30, 2001, where a candidate host galaxy of magnitude V$\\approx25$ has been found in the direction of the afterglow \\citep{fpk+02}. We note, however, that  the probability that this object is a foreground galaxy is not negligible ($P\\approx3\\times10^{-2}$, see e.g. eq.\\ref{eq:p2} in Sect.2.3). The {\\it high-redshift scenario} would naturally explain both the absence of optical counterpart and the high-energy spectrum, because the peak of the gamma-ray spectrum would be redshifted into the X-ray band.   If we are to use dark bursts to study obscured star formation in the universe \\citep{dfk+01}, we must first understand the source of the extinction. For those afterglows which are not at $z>5$  it is important to establish whether they are dark because of a dense circumburst medium \\citep{ry01}, or  their optical emission is extinguished by line-of-sight absorption from the medium of the host galaxy. We can use the properties of the afterglow, its location within the host galaxy, and the properties of the host galaxy itself to address this question. In this paper, we report observations of the burst \\grb\\ which was discovered and localized by {\\em BeppoSAX}.  A {\\em Chandra} observation of the {\\em BeppoSAX} error box enabled us to localize the likely host galaxy of the event and to identify a short-lived radio-transient, further refining the position to sub-arcsec accuracy.  From sensitive upper limits on the absence of an optical afterglow we estimate the amount of extinction by dust and from the X-ray spectrum the amount of absorbing gas. \\grb\\ appears to be the newest member of a small but growing group of well-localized dark bursts \\citep{fkb+99,tbf+00,dfk+01}. ",
        "conclusions": "In this paper we have presented the results of multi-wavelength observations of GRB\\,000210. This event was the brightest ever observed in $\\gamma$-rays in the {\\em BeppoSAX} GRBM and WFC. Nonetheless, no optical counterpart was found down to a limit of $R=23.5$.  GRB\\,000210 is therefore one of the events classified as dark GRBs, a class that makes up $\\approx 50\\%$ of all GRBs. It is still unclear whether  this class derives from a single origin or it is due to a combination of different causes. Some of these GRBs could be intrinsically faint events, but this fraction cannot be very high, because the majority of dark GRBs shows the presence of an X-ray afterglow similar to that observed in GRBs with optical afterglows \\citep{piro01,lcg01}. The most compelling hypotheses to explain the origin of dark bursts involve absorption, occurring either in the local environment of the GRB (circumburst or interstellar), or as Ly$\\alpha$ forest absorption for those bursts which have $z\\gsim5$.  As in the majority of bursts, \\grb\\ had an X-ray afterglow which was observed with {\\it BeppoSAX} and {\\it Chandra}.  The temporal behavior is well described by a power law, with a decay index $\\alpha_x=-1.38\\pm0.03$, similar to that observed in several other events  \\citep[e.g.][]{piro01}.  We did not find evidence for breaks in the light curve. The spectral index of the power law is also typical ($\\Gamma=1.95\\pm0.15$).  Thanks to the arcsecond localization provided by {\\em Chandra} we identified the likely host galaxy of this burst, determined its redshift ($z$=0.846) and detected a radio afterglow. The properties of the X-ray afterglow allowed us to determine the amount of dust obscuration required to make the optical afterglow undetectable ($A_R\\gtrsim2$). The X-ray spectrum shows significant evidence of absorption by neutral gas ($N_{HX}=(0.5\\pm0.1)\\times 10^{22}~\\ha$). However, we do not find evidence of a partially ionized absorber expected if the absorption takes place in a Giant Molecular Cloud, as recently suggested to explain the properties of the dark GRBs  \\citep[e.g.][]{ry01}. We conclude that, if the gas is local to the GRB, it has to be condensed in dense ($n\\gtrsim10^9$ cm$^{-3}$) clouds. We propose that these clouds represent the small-scale high-density fluctuations of the clumpy medium of star-forming GMCs.  Both the amount of dust required to extinguish the optical flux and the dust-to-gas ratio are consistent with those observed across the plane of our Galaxy. We cannot therefore exclude that the absorption takes place in the line-of-sight through  the interstellar medium of the host, rather than being produced by a GMC embedding the burst. This hypothesis is also consistent with the location of  GRB\\,000210 with respect to the center of the likely host galaxy. To explain the whole population of dark GRB this hypothesis would require that  host galaxies of GRBs should be characterized by quantities of dust and gas  much larger than typical,  arguing again for a physical connection between GRBs and star forming regions.   Finally, we discussed the possibility that the galaxy is unrelated to \\grb\\ and that it is a dark burst because it lies at $z\\gtrsim5$. In this case the X-ray absorbing medium should be substantially thicker than that observed in GRBs with optical afterglows.  Assuming that GRB\\,000210 is a typical representative of a population of events at high redshift, then these GRB's are embedded in a much denser environment than that of closer events.   Whichever of the explanations apply, it is clear that dark GRBs provide a powerful tool to probe their formation sites and possibly to explore the process of star formation in the Universe. We have at hand several observational tools to pursue this investigation. By increasing the number of arcsecond locations by radio, X-ray and far-infrared observations we can build up a sample of host galaxies of dark GRBs and study their distances and physical properties. The origin of the absorption in X-rays and optical can be addressed by broad band spectra and modeling, providing information on the dust and gas content of the absorbing structures. X-ray measurements are particularly promising in this respect for several reasons. First, X-rays do not suffer from absorption, in fact roughly the same number of dark GRB and GRBs with optical afterglows have an X-ray afterglow. Detection of X-ray lines can thus provide a direct measurement of the redshift. A comparative study of the X-ray properties of these two classes should also underline differences that can be linked to their origin, like the brightness of X-ray afterglows and the amount of X-ray absorption. Finally, measurements of variability of the X-ray absorbing gas would provide strong support to the {\\it local absorption} scenario. There are several mechanisms that can produce such a variability. The hard photon flux from the GRB and its afterglow will ionize the circumburst gas on short time scales, thus decreasing the effective optical depth with time \\citep{pl98}. The detection of a transient iron edge in GRB990705 \\citep{afv+00} and the decrease of the column density from the prompt to the afterglow phases in GRB980329 \\citep{fac+00} and GRB010222 \\citep{iz+01} are both consistent with this scenario. The variable size of the observable fireball that increases with the inverse of the bulk Lorentz factor can also produce variations of the column density, if the medium surrounding the source is not homogeneous. In this regard, two X-ray afterglows show some, admittedly marginal, evidence of variability of $N_H$ \\citep{pcf+99,yym+01}. In conclusion, the future of the investigations of {\\it dark} GRBs looks particularly bright."
    },
    "0201/gr-qc0201040_arXiv.txt": {
        "abstract": "Gravitational wave (GW) signals from coalescing binary neutron stars may soon become detectable by laser-interferometer detectors. Using our new post-Newtonian (PN) smoothed particle hydrodynamics (SPH) code, we have studied numerically the mergers of neutron star binaries with irrotational initial configurations. These are the most physically realistic initial conditions just prior to merger, since the neutron stars in these systems are expected to be spinning slowly at large separation, and the viscosity of neutron star matter is too small for tidal synchronization to be effective at small separation. However, the large shear that develops during the merger makes irrotational systems particularly difficult to study numerically in 3D. In addition, in PN gravity, accurate irrotational initial conditions are much more difficult to construct numerically than corotating initial conditions. Here we describe a new method for constructing numerically accurate initial conditions for irrotational binary systems with circular orbits in PN gravity. We then compute the 3D hydrodynamic evolution of these systems until the two stars have completely merged, and we determine the corresponding GW signals. We present results for systems with different binary mass ratios, and for neutron stars represented by polytropes with $\\Gamma=2$ or $\\Gamma=3$. Compared to mergers of corotating binaries, we find that irrotational binary mergers produce similar peak GW luminosities, but they shed almost no mass at all to large distances. The dependence of the GW signal on numerical resolution for calculations performed with $N\\gtrsim 10^5$ SPH particles is extremely weak, and we find excellent agreement between runs utilizing $N=10^5$ and $N=10^6$ SPH particles (the largest SPH calculation ever performed to study such irrotational binary mergers). We also compute GW energy spectra based on all calculations reported here and in our previous works.  We find that PN effects lead to clearly identifiable features in the GW energy spectrum of binary neutron star mergers, which may yield important information about the nuclear equation of state at extreme densities. ",
        "introduction": "\\label{sec:intro}  Coalescing neutron star (NS) binaries are likely to be one of the most important sources of gravitational radiation for the ground-based laser-interferometer detectors in LIGO \\cite{1}, VIRGO \\cite{2}, GEO600 \\cite{3}, and TAMA \\cite{4}. These interferometers are most sensitive to GW signals in the frequency range from about $10\\,{\\rm Hz}$ to $300\\,{\\rm Hz}$, which corresponds to the last several thousand orbits of the inspiral.  During this period, the binary orbit is decaying very slowly, with the separation $r(t)$ and phase $\\phi(t)$ following the standard theoretical treatment for the inspiral of two point masses (see, e.g., \\cite{Insp}). Theoretical templates for the corresponding quasi-periodic GW signals covering an appropriate range of values for the NS masses, as well as orbital phases and inclination angles, can be calculated to great precision (see, e.g., \\cite{Tem} and references therein), and template matching techniques can therefore be used to extract signals from noisy interferometer data. When the binary separation $r(t)$ has decreased all the way down to a few NS radii, the system becomes dynamically unstable \\cite{LRS} and the two stars merge hydrodynamically in $\\sim1\\,$ms. The characteristic GW frequency of the final burst-like signal is $\\gtrsim 1\\,{\\rm kHz}$, outside the range accessible by current broadband detectors. These GW signals may become detectable, however, by the use of signal recycling techniques, which provide increased sensitivity in a narrow frequency band \\cite{SR}. These techniques are now being tested at GEO600, and will be used by the next generation of ground-based interferometers. Point-mass inspiral templates break down during the final few orbits. Instead, 3D numerical hydrodynamic calculations are required to describe the binary merger phase and predict theoretically the GW signals that will carry information about the NS equation of state (EOS).  The first hydrodynamic calculations of binary NS mergers in Newtonian gravity were performed by Nakamura, Oohara and collaborators using a grid-based, Eulerian finite-difference code \\cite{ON}.  Rasio and Shapiro (\\cite{RS}, hereafter RS) later used Lagrangian SPH calculations to study both the stability properties of close NS binaries and the evolution of dynamically unstable systems to complete coalescence.  Since then, several groups have performed increasingly sophisticated calculations in Newtonian gravity, exploring the full parameter space of the problem with either SPH \\cite{Zhu,Dav,Ros} or the Eulerian, grid-based piecewise parabolic method (PPM) \\cite{New,Swe,Ruf}, and focusing on topics as diverse as the GW energy spectrum \\cite{Zhu}, the production of r-process elements \\cite{Ros}, and the neutrino emission as a possible trigger for gamma-ray bursts \\cite{Ruf}.  Some of these Newtonian calculations have included terms to model approximately the effects of the gravitational radiation reaction \\cite{Ruf}.  The first calculations to include the lowest-order (1PN) corrections to Newtonian gravity, as well as the lowest-order dissipative effects of the gravitational radiation reaction (2.5PN) were performed by Shibata, Oohara, and Nakamura, using an Eulerian grid-based method \\cite{Nak3}. More recently, the authors (\\cite{FR1,FR2}, hereafter Paper~1 and Paper~2, respectively), as well as Ayal et al.\\cite{Ayal}, have performed PN SPH calculations, using the PN hydrodynamics formalism developed by Blanchet, Damour, and Sch\\\"afer (\\cite{BDS}, hereafter BDS).  These calculations have revealed that the addition of 1PN terms can have a significant effect on the results of hydrodynamic merger calculations, and on the theoretical predictions for GW signals. For example, Paper~1 showed a comparison between two calculations for initially corotating, equal-mass binary NS systems. In the first calculation, radiation reaction effects were included, but no 1PN terms were used, whereas a complete set of both 1PN and 2.5PN terms were included in the second calculation.  It was found that the inclusion of 1PN terms affected the evolution of the system both prior to merger and during the merger itself.  The final inspiral rate of the PN binary just prior to merger was much more rapid, indicating that the orbit became dynamically unstable at a greater separation.  Additionally, the GW luminosity produced by the PN system showed a series of several peaks, absent from the Newtonian calculation.  Most previous hydrodynamic calculations of binary NS mergers have assumed corotating initial conditions, and many modeled the stars as initially spherical. However, real binary NS are unlikely to be described well by such initial conditions. Just prior to contact, tidal deformations can be quite large and the stars can have very nonspherical shapes \\cite{LRS}. Nevertheless, because of the very low viscosity of the NS fluid, the tidal synchronization timescale for coalescing NS binaries is expected to always be longer than the orbital decay timescale \\cite{BC}. Therefore, a corotating state is unphysical. In addition, at large separation, the NS in these systems are expected to be spinning slowly (see, e.g., \\cite{FK}; observed spin periods for radio pulsars in double NS systems are $\\gtrsim 50\\,$ms, much longer than their final orbital periods). Thus, the fluid in close NS binaries should remain nearly {\\it irrotational\\/} (in the inertial frame) and the stars in these systems can be described approximately by irrotational Riemann ellipsoids \\cite{LRS,BC,LomRS,IRRE}. In Paper~2, we showed that nonsynchronized initial conditions can lead to significant differences in the hydrodynamic evolution of coalescing binaries, especially in the amount of mass ejected as a result of the rotational instability that develops during the merger.  There are many difficulties associated with numerical calculations of binary mergers with irrotational initial conditions, especially in PN gravity.  Foremost of these is the difficulty in preparing the initial, quasi-equilibrium state of the binary system. In synchronized binaries, the two stars are at rest in a reference frame which corotates with the system, and thus relaxation techniques can be used to construct accurately the hydrostatic equilibrium initial state in this corotating frame (see \\cite{RS} and Paper~1). For irrotational systems, no such frame exists in which the entire fluid would appear to be in hydrostatic equilibrium. Instead, one must determine self-consistently the initial velocity field of the fluid in the inertial frame. Otherwise (e.g., when simple spherical models are used), spurious oscillations caused by initial deviations from equilibrium can lead to numerical errors. This is especially of concern in PN gravity, where the strength of the gravitational force at any point in the NS contains terms proportional to the gravitational potential and pressure at 1PN order.  Another serious problem for hydrodynamic calculations with irrotational initial conditions is the issue of spatial resolution and numerical convergence. As was first pointed out in RS2, in a frame corotating with the binary orbital motion, irrotational stars appear counterspinning so that, when they first make contact during the coalescence, a vortex sheet is formed along the interface. This tangential discontinuity is Kelvin-Helmholtz unstable at all wavelengths\\cite{Dra}. The sheet is expected to break into a turbulent boundary layer which propagates into the fluid and generates vorticity through dissipation on small scales. How well this can be handled by 3D numerical calculations with limited spatial resolution is unclear.  The outline of our paper is as follows. Section~II presents a summary of our numerical methods, including a brief description of our PN SPH code, and an explanation of the method used to construct irrotational initial conditions in PN gravity. Additionally, we give the parameters and assumptions for all new calculations discussed in this paper. Section~III presents our numerical results based on SPH calculations for several representative binary systems, all with irrotational initial configurations, but varying mass ratios and NS EOS. To test our numerical methods, we also study the effects of changing the initial binary separation and the numerical resolution. Section~IV presents GW energy spectra calculated from the runs in this and previous papers, as well as a discussion of how the measurement of spectral features could constrain the NS EOS. A summary of our PN results and possible directions for further research, including the possibility of fully relativistic SPH calculations, are presented in Section~V. ",
        "conclusions": "We have performed several large-scale SPH calculations of NS binary coalescence, assuming an irrotational initial condition for the binary system. Table~\\ref{table:runs} summarizes the relevant parameters of all runs performed, listing the adiabatic exponent $\\Gamma$, the mass ratio $q$, the initial separation $r_0$, and the number of SPH particles.  We continue to use the same nomenclature for our SPH runs introduced in Paper~2, although we present here a new run E1, using an improved initial configuration calculated by the method described in the previous section. Run E1 is for a system with a $\\Gamma=3$  EOS, equal-mass NS, and an initial separation $r_0=4.0$, and is similar in all respects to run B1 of Paper~2 (called the PN run in Paper~1), except that the initial condition is irrotational. It was continued until a quasi-stationary remnant configuration was reached. Run E2 is for a system with the same $\\Gamma=3$ EOS, but a mass ratio of $q=0.8$, and was started from a smaller initial separation $r_0=3.5$, since binaries with smaller masses take longer to coalesce. In runs F1 and F2, the EOS has $\\Gamma=2$, but we use the same mass ratios and initial separations as runs E1 and E2, respectively.  Additionally, to assess the numerical accuracy and convergence of calculations for irrotational binaries, we performed three runs which were primarily designed to test the dependence of the physical results on numerical parameters.  First, we performed one run, labeled T2, identical to E1 except for a smaller initial separation $r_0=3.5$.  This run was used to study how accurately the irrotational flow is maintained during the early stages of inspiral, and the effects of a small amount of spurious tidal synchronization on the GW signal. We then repeated run T2 using $5000$ and $500,000$ SPH particles per NS (for a total of $N=10^4$ and $N=10^6$ particles, respectively), to study the effect of numerical resolution on calculations where we know small-scale instabilities will develop.  The number of neighbors was adjusted in the two runs to be $N_N=50$ and $N_N=200$, respectively. This choice is dictated by the convergence and consistency properties of our basic SPH scheme: convergence toward a physically accurate solution is expected when {\\it both\\/} $N\\rightarrow\\infty$ {\\it and\\/} $N_N\\rightarrow\\infty$, but with $N_N/N\\rightarrow 0$ \\cite{RasThes}. The primary consideration behind the choice of the initial separation at $r_0=3.5$ (rather than $r_0=4.0$) was the high computational cost of a run with $N=10^6$.    Using the PN Lagrangian SPH code described in our previous papers, which is complete to 1PN order and includes radiation reaction effects, we have investigated a wide parameter space of binary NS mergers started from an irrotational initial condition. This initial configuration represents the most realistic approximation for NS binaries at separations of $r_0\\simeq 3.5-4.0\\,R$, corresponding to the onset of dynamical coalescence.  Based on these calculations, as well as those from our previous papers, we have calculated the energy spectrum of the gravitational wave emission for a variety of systems.  The key result is the existence of a ``cliff frequency'' in all of the energy spectra we have studied, i.e., a frequency above which the energy emitted dips dramatically beneath the point-mass approximation.  We attribute this effect to the onset of dynamical instability in binary systems, which leads to a much more rapid inspiral than the point-mass formula predicts, an effect amplified when PN terms are taken into account.  If the cliff frequencies of binary systems are even smaller than those found here (typically around $500\\,{\\rm Hz}$ for our standard NS parameters) when general relativity is treated consistently, they may lie within the frequency band accessible to broad-band laser interferometers. For example, proposals for LIGO II \\cite{LIGO2} place its upper frequency limit around $1000\\, {\\rm Hz}$ (where the sensitivity in terms of a characteristic GW strain has been degraded by a factor $\\sim 10$ due to photon shot noise), meaning that cliff frequencies should be seen if the true physical NS radius is sufficiently large, $R\\gtrsim 10\\, {\\rm km}$. Our results also suggest that the high frequency features in the GW energy spectrum, which result from emission during the merger itself ($f\\simeq 1600\\,{\\rm Hz}$) and from late-time oscillations of the remnant (\"ring down\" $f\\simeq 2200\\,{\\rm Hz}$), will be observable only by more advanced narrow-band detectors, but if observed, could place strong constraints on the NS EOS.  We also find from our calculations that initially irrotational binaries evolve in a qualitatively different way than do initially synchronized systems. Regardless of the choice of EOS, runs started from an irrotational configuration result in much less mass shedding than do synchronized runs, depositing no more than $1\\%$ of the total system mass in an outer halo which remains bound to the merger remnant at the center of the system. Additionally, there is a significant difference in the inspiral rate of such systems immediately prior to merger, with synchronized systems merging much more rapidly. Thus, since real NS binaries should be essentially irrotational, it is important to exercise caution when interpreting the results drawn from calculations for initially synchronized systems, both before and after the merger occurs. The lack of mass shedding seen in our PN calculations of irrotational systems leads us to conclude that essentially no mass should be shed from realistic binary systems, since general relativistic effects further suppress the mass shedding instability \\cite{Oech2,ShU}.  From a calculational standpoint, we find competing arguments for the ideal initial binary separation of an irrotational system. Runs started from larger separations (here $r_0=4.0R$) require greater computational resources, and show more spurious synchronization because of the numerical viscosity of the SPH method. Such calculations do provide a better treatment of the deviations from point-mass inspiral prior to merger which result from both Newtonian and PN finite-size effects. However, runs started from closer in ($r_0=3.5R$) are more reliable for drawing conclusions about the final state of the system, with regard to mass shedding as well as the rotational profile of the remnant.  We find, reassuringly, that the phase and amplitude evolution of the GW signal during the primary luminosity peak is unaffected by the initial separation.  We believe that our results are unaffected by the limited numerical resolution of 3D calculations. Although merging binary NS systems develop small-scale instabilities, whose evolution we cannot follow exactly, we see little effect on the GW signals we compute, so long as we use a sufficiently large number of SPH particles.  This is especially true for the phase of the GW signal. We conclude that numerical convergence for a given set of initial conditions and physical assumptions is possible without requiring excessive computational resources.  Based on these results, we believe that the fundamental limits of our method are not set by the numerical resolution or available computational resources, but rather by shortcomings in the PN formalism itself.  All our PN calculations are limited by the magnitude of the 1PN terms that appear in the hydrodynamic equations.  Since we cannot treat physically realistic NS models, we deal with the 1PN terms at reduced strength.  This hybrid method does in some sense approximate the cancellations found in GR between 1PN and higher-order terms, but in the end cannot fully model the non-linear nature of relativistic gravity.  To do so in a more complete manner, two different approaches have been employed.  The first is to attempt to integrate the hydrodynamics equations in full GR, as done recently by Shibata and Uryu \\cite{ShU}. Starting from quasi-equilibrium irrotational binary systems at the moment of first contact, they calculate the fully relativistic evolution of the system.  Such calculations represent a remarkable step forward, and can be considered the forefront of the current efforts to understand binary NS coalescence. They are limited only by the accuracy to which they can prepare their initial conditions and by the difficulty of extracting accurate GW signals from the boundary of 3D grids at a finite distance, usually well within the near zone of the source. As of yet no formalism has proven stable enough to handle a calculation which starts from the dynamically stable region and ends after the merger is complete.  A second approach is to use an approximation of full GR which is known to be numerically stable, known as the {\\it conformally flat\\/} (CF) approximation \\cite{Oech2,MMW,Bau}. In this approximation to general relativity, assuming a specific form for the spatial part of the metric allows the equations of GR to be reduced to a set of linked non-linear {\\it elliptic\\/} equations.  The drawback to the method is that the CF approximation is time-symmetric, and thus does not include the dissipative gravitational wave effects seen in full GR.  Such terms can be added externally to the formalism, though, to give the proper dynamical behavior \\cite{Oech2}. The CF approximation has been used successfully to compute quasi-equilibrium binary sequences to high accuracy for both synchronized \\cite{Phil,Phil2} and irrotational \\cite{Phil,Phil2,Irr} binaries, showing good agreement with fully relativistic calculations of synchronized \\cite{RelSyn} and irrotational binaries \\cite{RelIrr}, except for the case of extremely compact NS at small separations. Additionally, a PN variant of the CF approximation has been used to study rapidly rotating single-star configurations \\cite{DiffRot}, giving excellent agreement with fully GR calculations \\cite{GRDiffRot}. In addition to equilibrium studies, dynamical SPH calculations have recently been performed using relaxed, initially synchronized binary configurations \\cite{Oech2}.  The authors are currently working on a code that will take as an initial condition fully relativistic irrotational, quasi-equilibrium models for binary NS systems, calculated using spectral methods \\cite{Phil}.  The system will then be evolved in the CF approximation.  As the formalism is known to work for both binary configurations and rapidly rotating single-star configurations, reflective of our merger remnants, the hope is that such a method will allow us to calculate the evolution of a physically realistic binary system from the dynamically stable regime through merger and formation of a merger remnant, in a way that is consistent with GR throughout.  The comparison of such calculations with those performed in full GR will serve as an important check of the shortcomings and successes of either approach."
    },
    "0201/astro-ph0201231_arXiv.txt": {
        "abstract": "\\baselineskip=10pt Using adaptive optics on the W.M.\\ Keck II telescope we imaged Titan several times during 1999 to 2001 in narrowband near-infrared filters selected to probe Titan's stratosphere and upper troposphere.  We observed a bright feature around the south pole, possibly a collar of haze or clouds.  Further, we find that solar phase angle explains most of the observed east-west brightness asymmetry of Titan's atmosphere, although the data do not preclude the presence of a `morning fog' effect at small solar phase angle. ",
        "introduction": "\\begin{deluxetable}{ccccccc} \\tabletypesize{\\scriptsize} \\tablecaption{Parameters of Titan Observations. \\label{tbl-1}} \\tablewidth{0pt} \\tablehead{ \\colhead{} & \\colhead{} & \\colhead{} & \\colhead{} & \\colhead{} & \\colhead{Sub-Earth} & \\colhead{Solar} \\\\ \\colhead{UT} & \\colhead{UT} & \\colhead{Filter} & \\colhead{Exposure} & \\colhead{Apparent} & \\colhead{Latitude,} & \\colhead{Phase Angle,} \\\\ \\colhead{Date} & \\colhead{Time} & \\colhead{} & \\colhead{Time} & \\colhead{Size\\tablenotemark{a}} & \\colhead{Longitude\\tablenotemark{a,b}} & \\colhead{Position Angle\\tablenotemark{a,c}} \\\\ } \\startdata 30 October 1999 & 11:00 & J1158 & 3$\\times$120 sec & 0$\\farcs$865 & -19$\\fdg$92, 108$\\fdg$88 & 0$\\fdg$8870, 91$\\fdg$42 \\\\ 30 October 1999 & 11:36 & H1702 & 3$\\times$120 sec & 0$\\farcs$865 & -19$\\fdg$92, 109$\\fdg$45 & 0$\\fdg$8842, 91$\\fdg$48 \\\\ 17 August 2000 & 15:36 & J1158 & 4$\\times$120 sec & 0$\\farcs$774 & -24$\\fdg$05, 208$\\fdg$84 & 6$\\fdg$3375, 77$\\fdg$86 \\\\ 17 August 2000 & 15:13 & H1702 & 4$\\times$120 sec & 0$\\farcs$774 & -24$\\fdg$05, 208$\\fdg$48 & 6$\\fdg$3374, 77$\\fdg$85 \\\\ 20 February 2001 & 6:24 & J1158 & 4$\\times$120 sec & 0$\\farcs$774 & -23$\\fdg$14, 108$\\fdg$14 & 6$\\fdg$1893, 256$\\fdg$35 \\\\ 20 February 2001 & 5:50 & H1702 & 4$\\times$120 sec & 0$\\farcs$774 & -23$\\fdg$14, 107$\\fdg$60 & 6$\\fdg$1895, 256$\\fdg$35 \\\\  \\enddata   \\tablenotetext{a}{Obtained from the JPL Horizons Ephemeris, available at http://ssd.jpl.nasa.gov/horizons.html} \\tablenotetext{b}{Latitude and longitude coordinates are planetographic.} \\tablenotetext{c}{These coordinates describe the position of the sub-Solar point relative to the sub-Earth point on Titan.  The phase angle is defined as the Sun-Titan-Earth angle.  The position angle (PA) is defined CCW with respect to direction of the true-of-date Celestial North Pole.}   \\end{deluxetable}  Previous high-angular resolution studies of Titan have focused for the most part on mapping surface albedo \\citep{meier2000,gibbard1999,combes1997,smith1996,coustenis2001}. Several of these works derived atmospheric parameters, such as haze opacity, in order to separate the atmospheric and surface contributions.  Most recently, \\citet{coustenis2001} used moderate bandwidth filters ($\\sim0.15 \\mu$m FWHM) in the J- and H-bands on the 3.6 m Canada-France-Hawaii Telescope to image the surface and detected increased limb brightening around Titan's South Pole and on Titan's morning limb.  Titan's morning limb is to the East on the sky as viewed from Earth.  \\citet{coustenis2001}'s observations were just past Titan opposition from Earth with a Sun-Titan-Earth angle of only 0$\\fdg$5.  If, as one might expect, solar phase angle determines the East-West asymmetry in limb-brightening, then \\citet{coustenis2001} would have seen increased limb brightening on Titan's evening limb.  They interpreted this discrepancy from expectation by suggesting that a `morning fog' exists on Titan due to small diurnal changes in the thermal structure of the atmosphere and condensation of ethane as a parcel of air moves from the dark side of Titan to the sunlit side.  The data we present here show that at larger solar phase angles, solar phase angle dominates the East-West limb-brightening asymmetry; however, our results are consistent with a `morning fog' which would determine the East-West limb-brightening asymmetry at smaller phase angles. We also resolve the brightening seen at the southern limb by \\citet{coustenis2001} and \\citet{meier2000} into a `collar' around the south pole at a planetographic latitude of 70$\\fdg$S to 75$\\fdg$S.  Since October 1999 we have imaged Titan using the adaptive optics (AO) system on the W.M.\\ Keck II telescope \\citep{wizinowich2000} in several narrowband filters.  In this Note we present the data from narrowband filters centered at 1.158 $\\mu$m and 1.702 $\\mu$m, chosen to probe only the upper atmosphere of Titan above the middle-to-upper troposphere ($>\\sim 20$ km altitude).  The contribution from surface reflection is negligible; essentially all the light we observe in these filters is sunlight scattered by particles in the atmospheric haze layers \\citep{lemmon1995}. These images show a number of features in Titan's atmosphere: a bright collar around Titan's south pole, northern polar limb-brightening, and an east-west assymmetry in limb-brightening. ",
        "conclusions": ""
    },
    "0201/astro-ph0201188_arXiv.txt": {
        "abstract": "The association of G343.1-2.3 and PSR 1706-44 has been controversial from its first proposal. In this paper we address the difficulties, and argue that the association is still likely. New evidence comes from images of G343.1-2.3 obtained using the Australia Telescope Compact Array (ATCA), and the pulsar obtained using the {\\it CHANDRA} X-ray observatory. Mosaicing was required to cover the full extent of G343.1-2.3, and we present the polarisation images from this experiment. Also an X-ray pulsar wind nebula has been found in the archived {\\it CHANDRA} observations, with the correct morphology to support the association. The ATCA observations confirm the much larger extent of the SNR, which now encompasses the pulsar. The X-ray morphology points back toward the centre of the SNR, indicating the direction of the proper motion, and that the PSR and SNR are associated. ",
        "introduction": "The pulsar PSR B1706-44 is one of only seven that are known to emit gamma-rays in the GeV energy range \\citep{thompson_92} and one of only three that have been detected in the TeV range \\citep{kifune_95,chadwick_98}. \\citet{ros_1706} reported that the power law spectrum from the compact source at the site of the pulsar was best explained by synchrotron origin, presumably a nebula around the pulsar. \\citet{finley_98} found that PSR B1706-44 is surrounded by a compact X-ray nebula of radius $\\sim0.3$ pc using {\\it ASCA} data. This is also true of the other two TeV emitters, the Crab and Vela pulsars. The TeV gamma-rays, which appear to be unpulsed, can be explained as resulting from the Inverse Compton interaction of the high energy electrons, which produce the X-ray synchrotron emission, with ambient photons \\citep{aharonian_97}.   The original discovery that the {\\it COS-B} gamma-ray source was a pulsar was made by \\citet{PSR1706_discover}. The spin down luminosity is high, $3.4 \\times 10^{36}$~ergs, as would be expected. \\citet{1706} published a map made by the MOST telescope at 843~MHz of the area around the gamma-ray source. It showed a semicircular arc of emission, which has subsequently been denoted as the supernova remnant (SNR) G343.1-2.3, with the pulsar seemingly embedded in a small feature at its south eastern extremity. It was argued that the approximate distance of 3~kpc for the remnant, derived from the surface brightness-diameter relationship ($\\Sigma$-D) was compatible with the pulsar's dispersion measure, although the distance indicated by the widely used interstellar electron density model of \\citet{TC_model} would be only 1.8~kpc. The $\\Sigma$-D relationship was based on the flux density values from single dish observations, as both the VLA and the MOST integrated flux fell short of what would be expected, implying that the broad structure was being resolved \\citep{1706}.  If the association of the pulsar and the supernova remnant is real, a transverse velocity of $\\sim900$~kms$^{-1}$ is required for the pulsar to have moved from the approximate geometric centre of the SNR arc in the characteristic spin down time of 17~kyr. This is high, but not the highest measured value for a pulsar, nor as high as that implied for some seemingly well-established pulsar-SNR associations \\citep{duck_1,duck_2}.  \\citet{young_pulsars} imaged the area around the pulsar at 20~cm and 90~cm with the VLA and cast doubt on its association with G343.1-2.3. These arguments were based on the morphology, and the fact that the dispersion measure (DM) and $\\Sigma$-D distance were discrepant. The latter was subsequently countered by a 21~cm hydrogen line absorption measurement of the pulsar by \\citet{korbalski_HI} which gives a distance range of 2.4 to 3.2~kpc. Most recently the region has been imaged in more detail by \\citet{giacani_01} at 20, 6 and 3.6~cm and they concurred with \\citet{young_pulsars} description of a nebula size of $3.5^\\prime \\times 2.5^\\prime$.  \\citet{johnston_1706} also argue against the association. They report a maximum value of only 27~kms$^{-1}$ for the magnitude of the transverse velocity of the pulsar calculated from the interstellar scintillation. However in % their later paper \\citep{scint_49}, this estimate was revised up to 100~kms$^{-1}$. The possibility of a direct determination of the proper motion of PSR B1706-44 has been investigated, but no phase reference sufficiently strong for the current Australian VLBI network (the LBA) could be found.  We have made a number of high sensitivity, low surface brightness images of the remnant to determine its full extent. Furthermore we have used {\\it CHANDRA} archived data to search for a pulsar wind nebula (PWN) to provide evidence for the pulsar interaction with the interstellar medium (ISM). ",
        "conclusions": "Initially the arguments against the association were based on Kaspi's criteria \\citep{kaspi_assoc}. These can be summarised in order of ease of observability as; same position, same distance, same age, velocity reasonable and proper motion correct. It was felt that the pulsar DM distance and the SNR $\\Sigma$-D distance did not match, and the pulsar was at the edge of the SNR and had too low a velocity to have travelled from the geometric centre of the SNR. Also there was no morphological evidence to imply a velocity direction, and the scintillation velocity was too low. Since then work \\citep{korbalski_HI} has shown that the DM distance was incorrect, and our ATCA observations have found that the extent of the SNR has been underestimated. The Chandra data shows a morphological signature best interpreted as the path of the pulsar. The proper motion of the pulsar is the remaining question, and work is under way to attempt to answer that.  It is clear from both the low total flux values found by the MOST and VLA maps that there is a smooth broad component to this SNR. The ATCA observations confirm the extent of the broad emission, and that this is non-thermal in character. Taking this smooth component as the true limits of the SNR enlarges the extent and moves it southwards. This places the pulsar PSR1706-44 within the SNR shell, rather than on the rim, thus satisfying Kaspi's first criteria.  The X-ray PWN observed by {\\it CHANDRA} is consistent with the pulsar travelling from the general direction of the SNR centre, with either a very slow velocity (which would agree with the scintillation results) or a low local density (in agreement with the {\\it ROSAT} results).  We expect a SNR to be associated with PSR1706-44 because of its youth. The association with G343.1-2.3 was rejected in the past, but we have shown that the grounds for the rejection are not compelling. The question of the pulsar velocity could be resolved by finding a phase reference for VLBI observation, with either the SKA or an upgrade to the current LBA. We will be able to rule out the upper velocity limits with the ATCA within the next few years but will not be able to measure the lower limits for a few decades.  Finally the result stresses the importance of improving the ROSAT observations of this SNR with a mosaic over the whole area of the radio emission using either {\\it CHANDRA} or {\\it XMM}."
    },
    "0201/astro-ph0201141_arXiv.txt": {
        "abstract": "{ We report the detection of a thick disk in the edge-on, low surface brightness (LSB), late-type spiral \\gal, based on ultra-deep images in the $V$ and $R$ bands obtained with the VLT Test Camera during Science Verification on UT1. All steps in the reduction procedure are fully described, which, together with an extensive analysis of systematic and statistic uncertainties, has resulted in surface brightness photometry that is reliable for the detection of faint extended structure to a level of $V = 27.5~$ and $R = 28.5~$\\Msqarc.  The faint light apparent in these deep images is well-modeled by a thick exponential disk with an intrinsic scale height about 2.5 times that of the thin disk, and a comparable or somewhat larger scale length. Deprojection including the effects of inclination and convolution with the PSF allow us to estimate that the thick disk contributes 20-40\\% of the total (old) stellar disk luminosity of \\gal. To our knowledge, this is the first detection of a thick disk in an LSB galaxy, which are generally thought to be rather unevolved compared to higher surface brightness galaxies. ",
        "introduction": "Outside our own Galaxy, most of what we know about the structure, evolution and dynamics of stellar populations, and their connection to dark matter, is deduced from high surface brightness features: bars, bulges, and thin disks.  Fainter surface brightness components such as stellar halos, thick disks, and globular clusters probe galactic potentials differently, in both time and space owing to their larger age and extent.  The formation mechanisms of these faint tracers are still a matter of some controversy; suggestions range from early protogalactic collapse, secular processes such as heating from molecular clouds, black holes and spiral structure, through to later stochastic processes such as accretion (see recent reviews by \\cite{buser00}; \\cite{bhf00}; and references therein).  These scenarios predict different kinematical, morphological and chemical characteristics, but too few systems have been sufficiently well studied to constrain the models.  Due to the difficulty in detecting low surface brightness features reliably in external galaxies, the important complementary information they contain has only begun to be tapped.  In the Milky Way, faint disk and halo components can be separated on the basis of their kinematics and morphology, and -- to a certain extent -- metallicity, because individual stars can be resolved.  The Galactic stellar halo of field stars and the globular cluster systems have volume densities that decrease with galactocentric radius $r$ roughly as $\\rho (r) \\propto r^{-3.0}$ or $r^{-3.5}$ (\\cite{harris79}; \\cite{saha85}; \\cite{zinn85}), similar to results for halo populations in large spirals like M31 (\\cite{racine91}; \\cite{reitzel98}) and NGC~4565 (\\cite{fleming95}). Giant ellipticals and superluminous CD galaxies, on the other hand, which are thought to be the product of many mergers, have halo luminosities and globular cluster systems that fall less steeply, roughly as $\\rho (r) \\propto r^{-2.3}$ (\\cite{harris86}; \\cite{bridges91}; \\cite{harris95}; \\cite{graham96}).  The total mass, the bulk of which is believed to be contained in dark matter halos, is inferred from kinematical studies to have volume densities that decline as $\\rho (r) \\sim r^{-2}$ beyond a few disk scale lengths (see \\cite{sackett96} for a review).  Our Galaxy also has a faint thick disk whose density falls exponentially with increasing height ($z$) above the plane as $e^{-z/h_{z}^{\\rm thick}}$.  Its scale height $h_{z}^{\\rm thick} \\simeq 1 \\pm 0.3$~kpc (\\cite{reid93}; \\cite{ojha96}; \\cite{buser99}) is about three times larger than that of the much brighter thin disk.  The scale length of the thick disk is $h_{R}^{\\rm thick} \\simeq 3 \\pm 1.5$~kpc (\\cite{buser99}), similar to that of the Galactic thin disk.  Despite this, the thick disk contibutes only 2-9\\% of the total local stellar disk light (\\cite{reid93}; \\cite{ojha96}; \\cite{buser99}), and perhaps $\\sim$13\\% of the total disk luminosity of the Milky Way (\\cite{morrison94}).   For external galaxies, morphology determined through integrated surface brightness photometry is the only current method to detect and characterize faint galactic components.  Detections of extended light that are perhaps indicative of a thick disk component with $h_{z}^{\\rm thick} \\simeq 1-2$~kpc have been reported in a few external edge-on galaxies.  Early detections of extra-planar light in excess of that associated with a thin exponential disk  were limited to SO (\\cite{burstein79}) and early-type spirals with significant bulges (\\cite{vdkruit81a}; \\cite{vdkruit81b}; \\cite{wakamatsu84}; \\cite{bahcall85}; \\cite{shaw89}; \\cite{degrijs96}), \\cite{morrison97}).  leading to the supposition that thick disks were found in older stellar systems with significant central concentrations (\\cite{vdkruit81a}; \\cite{hamabe89}, \\cite{degrijs97}).  This hypothesis is consistent with the lack of a thick luminous component around the small, Scd spiral NGC~4244 in deep $R$-band observations reaching to $R = 27.5$\\Msqarc\\ (\\cite{fry99}), and in the bulgeless Sd edge-on NGC~7321 (\\cite{matthews99}). On the other hand, observations indicate that there are individual exceptions. Multiband photometry of the later-type Sc spiral NGC~6504 (\\cite{dokkum94}) revealed extended light interpreted as a weak thick disk with $h_{R}^{\\rm thick} \\simeq 2$~kpc.   Faint light high above the plane of the well-studied, late-type, edge-on spiral NGC~5907 has further complicated the picture of extra-planar light in small- or no-bulge spirals. First detected at heights of 3 to 6~kpc above the plane in deep $R$-band observations ({\\cite{morrison94}), this extended emission is intriguing because it is unlike any known thick disk or stellar component, having instead a morphology similar to that inferred for the dark matter halo distribution of NGC~5907 (\\cite{sackett94}). Other workers have confirmed the presence of the faint light in other bands (BVRIJK), and showed that the extended light is redder than the thin stellar disk. If the faint light is due to a thick disk, it is unlike any other, with a scale length that is at least twice that of its thin disk (\\cite{morrison99}). The stellar population responsible for this faint light remains highly controversial, ranging from normal or metal-rich populations with steep IMFs (\\cite{lequeux96}; \\cite{rudy97}; \\cite{jc98}), old, metal-rich accreted populations with normal IMF (\\cite{lequeux98}), or exceedingly metal-poor or giant-poor populations with few resolvable stars at the tip of the RGB (\\cite{zepf00}). The controversy remains because the full spectral energy distribution is apparently inconsistent with any single explanation (e.g. \\cite{zepf00}; \\cite{yost00}). \\footnote{The discovery of a faint, long, very narrow arc of light apparently associated with NGC~5907 (\\cite{shang98}) led Zheng et al. (1999) to suggest that the extended light in the galaxy might be an artifact due to confusion from the arc and foreground objects. The arc clearly contributes some light to some positions near the galaxy, but is too narrow and asymmetric to be the cause of the symmetric extended light detected by Morrison et al. (1994). Zheng et al. (1999) report that their photometry suffers from systematics at light levels fainter than $R = 27$\\Msqarc. (Due to a large pixel size, the PSF was often undersampled, despite the seeing of 3.4 to 5.4\\arcsec\\ that was typical of their observations). Since all detections of faint extended light in NGC~5907 have been reported for $R \\geq 27$\\Msqarc, and all optical photometry (including that of \\cite{zheng99}) agrees above this level, the mystery of this faint halo light remains.}  The puzzling nature of the extended light in NGC~5907 has motivated new studies to test a possible connection between faint optical and IR light and dark matter in this and other spirals (\\cite{gilmore98}; \\cite{rauscher98}; \\cite{uemizu98}; \\cite{abe99}; \\cite{beichman99}; \\cite{yost00}; \\cite{zepf00}). The optical results are mixed, but infrared surface brightness photometry indicates that whatever produces the faint optical light detected to date does not appear to emit strongly at IR wavelengths far from the plane of the galactic disks. Thus, if associated with known stellar populations, the sources of the faint light are unlikely to account for the dark mass of spiral galaxies.   \\begin{table*}[t] \\begin{center}  \\caption{Basic Properties of \\gal} \\vspace*{0.5cm}  \\begin{tabular}{c c c} \\hline \\noalign{\\smallskip} \\multicolumn{1}{c}{Parameter} &\\multicolumn{1}{c}{Value} &\\multicolumn{1}{c}{Reference}\\\\ \\hline\\hline  $\\alpha$,$\\delta$ (J2000.0)\t&21 12 10.8, $-$37 37 38\t&\\cite{karach99} \\\\ type\t\t\t\t&Sc+6\t\t\t&\\cite{mathewson96} \\\\ redshift\t\t\t&7680$\\pm$ 10\\kms\t&\\cite{mathewson96} \\\\ inclination\t\t\t&88\\degr\t\t\t& this paper \\\\ PA\t\t\t\t&120\\fdg 4 $\\pm$ 0\\fdg 5\t& this paper \\\\ major-axis D$^{\\ast}$\t\t&86\\arcsec\\\t\t\t& this paper \\\\ m$_B$\t\t\t\t& 16.67$\\pm$0.09\t&\\cite{lauberts89}  \\\\ m$_V$\t\t\t\t& 16.40$\\pm$0.03\t\t& this paper \\\\ m$_R$\t\t\t\t& 15.92$\\pm$0.04\t\t& this paper \\\\ m$_I$\t\t\t\t& 15.47$\\pm$0.06\t&\\cite{mathewson96} \\\\ M$_R$\t\t\t\t& $-$19.1$\\pm$0.3\t\t\t& this paper \\\\ M$_V$\t\t\t\t& $-$18.7$\\pm$0.3\t\t\t& this paper \\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\end{center}  \\footnotesize  \\noindent{ } \\hspace*{3.4cm} $^{\\ast}$major-axis diameter measured from the R=27.0 \\Msqarc\\ contour.\\\\ \\hspace*{3.8cm} Magnitudes are not corrected for extinction.  \\label{tab:basic_props} \\end{table*}  In this paper, we report on the collection, reduction and analysis of ultra-deep surface photometry of the isolated, edge-on, low surface brightness, Sd galaxy \\gal, using some of the first science observations taken with the VLT. The simple optics, good seeing, and extremely well-sampled PSF of our observations ensured a low and well-understood level of scattered light and accurate identification of contaminating sources. Concurrent deep observations of unrelated blank fields with the VLT were used to create dark sky flat fields at the appropriate wavelengths. Considering all sources of uncertainty, including those from light scattered through the wings of the PSF, we conclude that the resulting surface photometry is reliable to a level of $R = 28.5$\\Msqarc\\ and $V = 27.5$\\Msqarc. Analysis of these data reveals a faint component that we interpret as a thick disk, to our knowledge the first thick disk discovered in an LSB galaxy.  In Sect.~2 we describe the VLT observations and observing strategy. In Sect.~3 the data reduction process, including the production of dark sky flats and the procedures for masking, mosaicing, calibrating, and determining the sky flux are outlined. The procedure to extract profiles from the deep images is given in Sect.~4, along with a brief description of the error analysis, which is discussed in depth in the appendix. The resulting $V$ and $R$ surface photometry of \\gal\\ are presented in Sect.~5, along with a description of the fitting procedure for the thin and thick disk parameters. A thorough analysis of scattered light due to the tightly-constrained PSF is discussed in Sect.~5, and ruled out as the cause of the faint extended light we detect in \\gal. The thin and thick disks, including their inferred intrinsic properties are described in Sect.~6.  We summarize and conclude in Sect.~7.  Throughout this paper we assume a distance of 102\\,Mpc to \\gal\\ (based on a Hubble constant of  H$_\\circ$=75\\,km/s/Mpc), which yields an image scale of 0.495\\,kpc per arcsecond. ",
        "conclusions": "We have used the VLT test camera on UT1 to obtain deep surface brightness photometry of the edge-on LSB galaxy \\gal\\ in the $V$ and $R$-bands.  Careful masking of foreground and background objects to obtain an accurate value of the sky flux on our science mosaics, and an analysis of flat-fielding uncertainties -- both statistical and systematic -- on a variety of spatial scales, allow us to estimate confidently the total uncertainty in our deep surface photometry.  We conclude that on the size scales important for probing faint, extended structure, we reach $V = 28$ and $R = 29~$\\Msqarc.  A detailed analysis of the PSF of the images, derived from faint isolated stars on the mosaic and standard stars, indicates that scattered light affects the extended vertical luminosity profiles of \\gal\\ only for $R > 28.5~$\\Msqarc.  Extended light in excess of that expected for a single-component thin disk is detected at about $R > 26.5~$\\Msqarc\\ in nearly all vertical profiles perpendicular to and up to 17~kpc along the major axis of \\gal. The same component may have also been detected in the $V$ band frames, but the lower $S/N$ of these frames and the larger PSF in $V$ make this detection less robust.   Given the geometric form of the extended light in this apparently bulgeless galaxy, we interpret the faint $R$-band light as a thick disk.  Two-component exponential disk fits were made to the observed surface brightness profiles and used to determine projected and -- after deprojection and deconvolution -- intrinsic structure parameters for the thin disk of \\gal\\ in the $V$- and $R$-bands and for the thick component in $R$. In particular, we find:  \\begin{itemize}  \\item The thin disk has projected scale heights perpendicular to the major axis of $h_z^{\\rm thin} = 380 \\pm 35~$pc in the $R$-band and $h_z^{\\rm thin} = 380 \\pm 45~$pc in $V$. The projected scale length of the thin disk is $8.9 \\pm 1.5~$kpc in both bands.  \\item   After deprojection and deconvolution with the PSF derived from our observations, we estimate that, within the errors, the intrinsic thin disk scale heights are 80\\% of the measured values.  The intrinsic scale length of the thin disk is $2/3$ of the fitted value.  \\item   The face-on central surface brightness of the thin disk is estimated to be  $\\mu_{0}^{\\rm thin} = 23.6~$\\Msqarc\\ in $R$ and $\\mu_{0}^{\\rm thin} = 24.1~$\\Msqarc\\ in $V$.  \\item   The thick disk, which is detected robustly in our $R$ surface photometry, has a projected scale height of $h_z^{\\rm thick} = 810 \\pm 40~$pc in the $R$-band and $h_z^{\\rm thick} = 760 \\pm 75~$pc in $V$. The projected scale length of the thick disk cannot be determined precisely from our observations, but within uncertainties is consistent with that of the thin disk.  \\item   The intrinsic scale parameters of the thick disk are somewhat smaller than these measured projected ones, though not as dramatically different as the projected and intrinsic parameters of the thin disk.  \\item   Simple considerations lead to an estimate for the face-on central surface brightness of the thick disk of  $24.1 < \\mu_{R,0}^{\\rm thick} < 24.9~$\\Msqarc.  \\item   The thick disk is likely to contribute 20-40\\% of the total (old) disk $R$-band luminosity of the galaxy.  \\item   The total central surface brightness of the (thin + thick) disk is $\\mu_{R,0} > 23.1~$\\Msqarc, which places \\gal\\ securely in the category of low surface brightness (LSB) galaxies.  \\end{itemize}  This detection of a thick disk adds to only a few others known in external galaxies (see Sect.~1), and to our knowledge is the first known thick disk in an LSB galaxy.  The thick and thin disks of \\gal\\ have similar scale heights as their corresponding components in the Milky Way, but larger scale lengths. Importantly, the thick disk of \\gal\\ appears to contribute a larger fraction of the overall old disk light than does the Galactic thick disk. (Young HII regions have been masked and so do not enter into the extrapolated estimates we have made.) A prominent thick disk in \\gal\\ is particularly interesting since, compared to their high surface brightness cousins, LSB galaxies are thought to be more dark-matter dominated and to have less evolved disks. The VLT observations reported here suggest that, at least in the case of \\gal, such an unevolved thin disk can coexist with a substantial thick luminous component, perhaps providing a clue to the formation mechanism of thick disks in all spirals."
    },
    "0201/astro-ph0201377_arXiv.txt": {
        "abstract": "We present the first edition of a catalog of variable stars from OGLE-II Galactic Bulge data covering 3 years: 1997--1999. Typically 200--300 $I$ band data points are available in 49 fields between $-11$ and $11$ degrees in galactic longitude, totaling roughly $11$ square degrees in sky coverage. Photometry was obtained using the Difference Image Analysis (DIA) software and tied to the OGLE data base with the DoPhot package. The present version of the catalog comprises 221,801 light curves. In this preliminary work the level of contamination by spurious detections is still about 10\\%. Parts of the catalog have only crude calibration, insufficient for distance determinations. The next, fully calibrated, edition will include the data collected in year 2000. The data is accessible via FTP. Due to the data volume, we also distribute DAT tapes upon request. ",
        "introduction": "\\label{sec:intro}  The main goal of the Optical Gravitational Lensing Experiment (OGLE, Udalski, Kubiak \\& Szyma\\'nski 1997) is to search for microlensing events. Observationally, these events are basically a rare type of an optical variable, therefore it came as no surprise that after several years microlensing experiments have an exceptional record of variability in terms of the number of objects and epochs. To maximize event rates, microlensing searches focus on monitoring of very crowded, and scientifically attractive, stellar fields; the Galactic Bulge region and Magellanic Clouds. Some observations are conducted in denser portions of the Galactic disk.  It is a common situation nowadays that the ability to generate data far exceeds the ability to process it, and even more so, to comprehend it. The list of projects which aim at monitoring significant parts of the sky for variability includes more than 30 names ({\\it http://www.astro.princeton.edu/faculty/bp.html}), yet only a small fraction of those can process the data efficiently enough to make the measurements publicly available soon after the data is taken (e.g. Brunner et al. 2001). The issue of exporting the data in a convenient form compounds the problem. The National Virtual Observatory (NVO) project has very ambitious plans to provide the tools and some standards (perhaps de facto standards) for processing the large amounts of information and web data publication (Szalay 2001). Large catalogs have added complexity (project description {\\it http://www.us-vo.org/}). By the time some sort of processing is complete, new information emerges in the process, frequently information which should be incorporated into the catalog. It seems that the only static layer is the raw data itself, typically CCD images, however the photometric output from number crunchers should also be reasonably slow to change with the new developments.  A regular practice in OGLE is to release the data in the public domain as soon as possible. The most significant contributions are: BVI maps of dense stellar regions (Udalski et al. 1998b, 2000a), Cepheids in Magellanic Clouds (Udalski et al. 1999a, 1999b), eclipsing variables in the SMC (Udalski et al. 1998a), catalogs of microlensing events (Udalski et al. 2000b, Wo\\'zniak et al. 2001). Examples from other microlensing teams include samples of MACHO microlensing events (Alcock et al. 1997a, 1997c) and selected variable star work from MACHO (1997b) and EROS (Afonso et al. 1999). In real time detection of microlensing events the main benefit comes from the follow-up work (e.g. Sackett 2000), in practice only possible with immediate publication on the WWW. Therefore, all major microlensing teams (OGLE, EROS, MACHO and MOA) have, or had, active alert systems.   \\noindent A recent contribution to the publicly available data on variable stars is a WWW interface to the MACHO database (Allsman \\& Axelrod 2001), which started with somewhat limited features, but has plans for expansion. Similar ideas of making evolving catalogs have been discussed within OGLE for some time now and are motivated by the challenges of data processing/accessibility. The main objective here is not to make a potential broad user wait for a long time until the team makes the final refined product. There is a lot of potential use from the data at all levels of processing, as demonstrated by the serendipitous recovery of high proper motion stars (Eyer \\& Wo\\'zniak 2001) and discovery of the longest microlensing event ever observed, most likely caused by a black hole with the mass of several solar masses (Mao et al. 2002), both found in preliminary OGLE catalogs. OGLE has just released an online catalog of $\\sim$70,000 candidate variables in the LMC and SMC (\\.Zebrun, Soszy\\'nski et al. 2001). With this paper we release an initial catalog of 221,801 candidate variables in the Galactic Bulge from Difference Image Analysis of OGLE-II data from seasons 1997-1999. Parts of the current edition are still not fully calibrated and should not be used in distance estimates (Section~\\ref{sec:phot}).  We restate the basic information about the data in Section~\\ref{sec:data} and in Section~\\ref{sec:phot} we briefly summarize the process of finding variability. Section~\\ref{sec:ftp} gives the details of how the catalog is structured, followed by final remarks and future plans in Section~\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  As mentioned before, the current edition of the catalog basically includes entire output of the DIA pipeline as described by Wo\\'zniak (2000), supplemented with determinations of the reference flux to put the light curves on the magnitude scale. Some optimization has been performed to keep the contamination by artifacts low without rejecting too many real variable objects, but it must be clearly stated that about 10\\% of the light curves in the present release are not real objects and result from various problems, undetected at the pipeline level. We are extending the work of Mizerski \\& Bejger (2001) from the first BUL\\_SC1 field to all fields in the effort to flag several common types of artifacts and clean the sample of spurious objects.  Classifying a real variable star versus a spurious one is the first step in the interpretation of this data. Ultimately we envision increasingly refined information added to the catalog to facilitate applications. This should include the full classification of the detected variables, cross identification with objects found by 2MASS and in X-ray catalogs, periods for periodic sources etc. The work on automated classification of periodic variables is in progress. In addition to examining 2-D projections of a multidimensional parameter space and trying to code a human made algorithm (see Mizerski \\& Bejger 2001 and Wo\\'zniak et al. 2001 for such work on this data), we are experimenting with data mining techniques. Even with the current volume of data in OGLE-II we believe it is enabling to make the transition from ''telling the computer how to do it'' to ''telling the computer what to do'' and leaving the rest to the algorithm. A number of standard machine learning tools are available, which take small preclassified subsets of light curves and can ''learn'' to classify the rest of the data.  \\.Zebrun, Soszy\\'nski et al. (2001) provided a convenient web interface to access the data on variables in the LMC and SMC. It is our intention to build a similar tool with the addition of positional searches. Although the volume of the data which can be accessed by browsing web pages is limited in practice, the search by coordinates is a powerful tool for numerous applications. Transfer by FTP and distribution of DAT tapes are currently primary modes of accessing major parts of this catalog. For your copy of a DAT tape, please contact Prof. Bohdan Paczy\\'nski (email: bp@astro.princeton.edu, mail: Princeton University Observatory, Princeton, NJ, 08544). To access this archive online use the OGLE web site {\\it http://bulge.princeton.edu/$\\sim$ogle/ogle2/bulge\\_dia\\_variables} ."
    },
    "0201/astro-ph0201363_arXiv.txt": {
        "abstract": "{Using two epochs of HST/WFPC2 images of the metal-rich globular cluster NGC 6528 we derive the proper motions of the  stars and use them to separate the stars belonging to NGC 6528 from those of the Galactic bulge. The stellar sequences in the  resulting colour-magnitude diagram for the cluster are significantly better determined than in previously published data. From comparison of the colour-magnitude diagram with the fiducial line for NGC 6553  from Zoccali et al.  (2001) we conclude that the two globular clusters have the same age. Further, using $\\alpha$-enhanced stellar isochrones, NGC 6528 is found to have an age of $11\\pm 2$ Gyr. This fitting of isochrones also give that the cluster is 7.2 kpc away from us. From the measured velocities both the proper motion of the cluster and the velocity dispersion in the Galactic bulge are found.  NGC 6528 is found to have a proper motion relative to the Galactic bulge of $<\\mu_l>$ = 0.006 and $<\\mu_{\\rm b}>$ = 0.044 arcsec per century. Using stars with $\\sim 14 <V_{\\rm 555} < 19$ (i.e. the red giant branch and horizontal branch) we find, for the Galactic bulge,  $\\sigma_{\\rm l}= 0.33\\pm 0.03$  and $\\sigma_{b}= 0.25\\pm0.02$ arcsec per century. This give $\\sigma_{\\rm l}/\\sigma_{\\rm b}=1.32\\pm0.16$,  consistent both with previous proper motion studies of K giants in the  Galactic bulge as well as with predictions by models of the kinematics of bulge stars. ",
        "introduction": "NGC 6528 is perhaps {\\sl the} most metal-rich globular cluster known and several studies have therefore targeted this cluster, e.g. Ortolani et al. (1992), Ortolani et al. (1995), Richtler et al. (1998), Cohen \\& Sleeper (1995), Heitch \\& Richtler (1999), Carretta et al. (2001). It has also been used as a reference in studies of other clusters, e.g.  Davidge (2000).   We summarize important literature measurements of NGC6528 in Table \\ref{ngc6528.tab}.  NGC 6528 is at $(l,b)=(1.14,-4.12)$, i.e. in the plane of the Galactic disk and towards the Galactic bulge.  The first effect of this is that it is heavily reddened by foreground dust, Ortolani et al. (1992), Richtler et al. (1998), Heitsch \\& Richtler (1999). Most recent distance estimates put the cluster within less than 1 kpc from the Galactic centre (e.g. Richtler et al. 1998 and references therein). The close proximity to bulge stars further complicates the interpretation of the colour-magnitude diagram. Since the bulge stars and the cluster stars have roughly the same distance modulus they are superimposed in the colour-magnitude diagram. This effect has been noted to be particularly pronounced in the red giant branch (Richtler et al. 1998).  \\begin{table*} \\caption[]{Data for NGC6528 compiled from the literature.} \\label{ngc6528.tab} \\begin{tabular}{llllllllllllllllll} \\hline\\noalign{\\smallskip} & Value & Ref. & Comment\\\\\\noalign{\\smallskip} \\hline\\noalign{\\smallskip} Core radius & $0\\farcm 09$& Harris (1996) \\\\ Distance & 7.5 kpc  & Ortolani et al.\\,(1992) \\\\ $\\Delta(m-M)$ & 16.4 & Zinn (1980) & 19.1 kpc\\\\ & 14.35 &  Ferraro et al. (1999) \\\\ & 15.15$\\pm$0.24&  Heitsch \\& Richtler (1999) & isochrone fitting \\\\ {[Fe/H]}     & $+0.01$ & Zinn (1980)\\\\ & $+0.29$ & Bica \\& Patoriza (1983)\\\\ & $+0.12$ & Zinn \\& West (1984) \\\\ & $-0.23$ & Armandroff \\& Zinn (1988) & Integrated spectra IR Ca { \\sc ii}\\\\ & high, sim to NGC6553& Ortolani et al.\\,(1992)\\\\ & $-0.23$ &Origlia et al.\\,(1997) & IR abs. at 1.6 $\\mu$m \\\\ &  $+0.07\\pm0.1?$ & Carretta et a. (2001) \\\\ {[M/H]} & $+0.1/-0.4$ & Richtler et al.\\,(1998)&Trippico isochrone/Bertelli isochrone\\\\ & -0.31 &  Ferraro et al. (1999) \\\\ & 0.00 & Heitsch \\& Richtler (1999)          \\\\ $Z$ & $Z_{\\odot}$ & Bruzual et al.\\,(1997) & \\\\ Age & 14 Gyr & Ortolani et al.\\,(1992) & metallicity comparable to solar\\\\ & $12\\pm2$ Gyr & Bruzual et al.\\,(1997) & \\\\ E(B-V) &  0.56 & Zinn (1980) \\\\ &  0.55 & Ortolani et al.\\,(1992) & NGC6553 as reference and $\\Delta(m-M)_V= 14.39$\\\\ & 0.62 & Bruzual et al.\\,(1997) & \\\\ & 0.62 & Ferraro et al. (1999) \\\\ E(V-I)& 0.8/0.6 & Richtler et al.\\,(1998)&Trippico isochrone/Bertelli isochrone\\\\ & 0.46$\\pm$0.03&   Heitsch \\& Richtler (1999)           & isochrone fitting\\\\ \\noalign{\\smallskip} \\hline\\noalign{\\smallskip} \\end{tabular} \\end{table*}  In a pre-study we noted in particular that the colour-magnitude diagram of the globular cluster NGC 6528 closely resembled that of NGC 5927 {\\sl if} the latter was superimposed on the colour-magnitude diagram of Baade's window (for a set of representative colour-magnitude diagrams see Feltzing \\& Gilmore 2000).  This is consistent with the cluster being virtually inside the Galactic bulge and thus having a large contribution of bulge stars in its colour-magnitude diagram. Moreover, if the cluster is as metal-rich as indicated in previous studies (and now confirmed by Carretta et al. 2001) then the red-giant branch as well as both the turn-off and the horizontal branch of the bulge and globular cluster will appear at virtually the same magnitudes and colours. The bulge stars will be more spread out in the colour-magnitude diagram than those in the cluster, due to the large range of metallicities and ages present in the bulge (e.g.  McWilliam \\& Rich 1994, Feltzing \\& Gilmore 2000). Thus the only way to obtain a clean colour-magnitude diagram for the cluster is to obtain proper motions of the cluster stars relative to the bulge stars and separate the two populations using their proper motions.  This conclusion prompted us to apply for HST time to obtain a  second epoch of observations of NGC 6528 with WFPC2 in order to derive the relative proper motion of the cluster as compared to that of the Galactic bulge.  We report here on the results from this proper motion study. The article is organized as follows; Sect. 2 presents the data, derivation of photometry and proper motions are discussed in Sect. 3 and 4, Sect. 5 and 6.  contains an extensive discussion of mean proper motions and velocity dispersions for  NGC 6528 relative  to the Galactic bulge as well as the velocity dispersion in the bulge itself, in Sect. 7 we use the proper motion information to obtain  a clean colour-magnitude diagram for NGC 6528, Sect. 8 contains  a discussion of the age for the cluster, and finally Sect. 9  provides a brief discussion and conclusions. ",
        "conclusions": "Colour-magnitude diagrams of  the metal-rich globular cluster NGC 6528 are notoriously difficult to analyze. This is due to the fact the cluster is situated in the Galactic bulge and thus the fields stars belonging to the bulge have the same magnitudes as the stars in NGC 6528.  Using two epochs of observations with HST/WFPC2 we obtain the stellar proper motions for all stars in the field. The proper motion are  used to separate the bulge from the cluster stars. The stellar sequences in the resulting colour-magnitude diagram are better defined than in any previously published colour-magnitude diagram.  Using $\\alpha$-enhanced stellar isochrones we find NGC 6528 to have a probable age of $11\\pm 2$ Gyrs, this is the first attempt to establish the absolute age of NGC 6528. Previous studies have only compared the fiducial ridge line for the cluster to that of other globular clusters of similar metallicities.  Mainly the comparisons have been with regards to NGC 6553 and 47 Tuc.  With the new metallicity determinations for individual stars in both NGC 6553 and NGC 6528 it is now clear that 47 Tuc  (at $-0.71$ dex) has a  significantly lower metallicity than NGC 6528 and NGC 6553 and is thus not a suitable comparison cluster as regards differential age determinations (see e.g. Stetson et al. 1996 and references therein).  The distance modulus obtained when fitting the stellar isochrones yields a distance to NGC 6528 of 7.2 kpc, fully compatible with previous derivations.  A comparison with the fiducial line for NGC 6553 confirms results in earlier studies, e.g. Ortolani et al. (1995), that the two clusters indeed have very similar ages.  The stellar proper motion also provide velocity dispersions for both the cluster and field stars. The velocity dispersion of the cluster is most likely dominated by measurement errors. The bulge dispersion can thus be found by deconvolution. The resulting dispersions are consistent with what has previously (Spaenhauer et al. 1992) been found for Bulge giants. Moreover, combining our results with those by Zoccali et al.  (2001) we are able to confirm the difference in $\\sigma_{\\rm l}$ at two  positions in the bulge as predicted by the model in Zhao (1996).  To our knowledge, our and Zoccali et al.'s study are the first to address the proper motions amongst bulge horizontal branch and fainter stars.  Finally we derive the mean proper motion of NGC 6528 relative to  the Galactic bulge and also the space velocities. We find  $<\\mu_{\\rm l}>=\\sim +0.006$ and $<\\mu_{\\rm b}>=\\sim +0.044$ arcsec per century and ($\\Pi$,$\\Theta$,$W$)=(-142, 303, 4)  km s$^{-1}$.  \\acknowledgement{We thank the Royal Swedish Academy of Sciences for a collaborative grant that enabled visits to Cambridge, for SF, and to Lund, for RAJ }"
    },
    "0201/astro-ph0201013_arXiv.txt": {
        "abstract": "We have derived the uncertainties to be expected in the derivation of galaxy physical properties (star formation history, age, metallicity, reddening) when comparing broad-band photometry to the predictions of evolutionary synthesis models. We have obtained synthetic colors for a large sample ($\\sim$9000) of artificial galaxies assuming different star formation histories, ages, metallicities, reddening values, and redshifts. The colors derived have been perturbed by adopting different observing errors, and compared back to the evolutionary synthesis models grouped in different sets. The comparison has been performed using a combination of Monte Carlo simulations, a Maximum Likelihood Estimator and Principal Component Analysis. After comparing the input and derived output values we have been able to compute the uncertainties and covariant degeneracies between the galaxy physical properties as function of (1) the set of observables available, (2) the observing errors, and (3) the galaxy properties themselves. In this work we have considered different sets of observables, some of them including the standard Johnson/Cousins ($UBVR_{C}I_{C}$) and {\\it Sloan Digital Sky Survey (SDSS)} bands in the optical, the {\\it 2 Micron All Sky Survey (2MASS)} bands in the near-infrared, and the {\\it Galaxy Evolution Explorer (GALEX)} bands in the UV, at three different redshifts, $z$=0.0, 0.7, and 1.4. This study is intended to represent a basic tool for the design of future projects on galaxy evolution, allowing an estimate of the optimal band-pass combinations and signal-to-noise ratios required for a given scientific objective. ",
        "introduction": "\\label{sec4}  Once these quantities had been derived we computed the mean differences between the output and input values along with the mean 1-sigma spread, at fixed intervals in the input properties. The bins used were 0.5\\,Gyr, 0.025\\,dex, and 0.025\\,mag in the formation timescale, age, and reddening, respectively. Mean differences and 1-sigma values in the stellar metallicity were computed for each of the input values considered.  In Figures~\\ref{fig3}a \\& b we show the results obtained before and after computing the mean differences and 1-sigma errors for a subsample of 500 nearby galaxies with errors in the colors of 0.07\\,mag and U+BVRI+JHK data available. Due to the relevance of the $K$-band luminosities in order to derive stellar masses in nearby (Arag\\'on-Salamanca et al$.$ 1993; GIL00a) and intermediate-redshift galaxies (see Brinchmann \\& Ellis 2000 and references therein), the mean differences between the derived and input $K$-band mass-to-light ratios were also computed. Mean uncertainties for all the sets of observables, redshifts, and observing errors considered are summarized in Table~\\ref{table3}.  In addition, we studied the degeneracies between the galaxy physical properties analyzing the distribution of the unitary PCA1 vector components. In Figure~\\ref{fig4}a we show the frequency histograms obtained for the sample of high-redshift galaxies assuming an error in the colors of 0.07\\,mag and the SDSS, SDSS+2MASS, and GALEX+SDSS+2MASS sets. Note that the PCA1 vector points toward the direction where the largest fraction of the galaxy properties' variance occurs. In this sense, a PCA1 vector with components ($u_{\\mathrm{log}~t}$,$u_{E(B-V)}$,$u_{\\mathrm{log}~Z/Z_{\\odot}}$,$u_{\\tau}$)=($+$0.707,$-$0.707,0,0), say, implies the existence of a degeneracy between age and reddening in the sense that younger, obscured stellar populations have colors that are indistinguishable from older but less extincted ones. In this case, no age-metallicity or age-timescale degeneracies would be present. However, the behavior described above could also result in a PCA1 vector with components ($u_{\\mathrm{log}~t}$,$u_{E(B-V)}$,$u_{\\mathrm{log}~Z/Z_{\\odot}}$,$u_{\\tau}$)=($-$0.707,$+$0.707,0,0). That, however, would appear in a different position at the frequency histogram shown in Figure~\\ref{fig4}a. Thus, in order to reduce this sign ambiguity when interpreting our results we forced the $u_{\\mathrm{log}~t}$ component to be positive, changing the sign of all the vector components if $u_{\\mathrm{log}~t}$ was negative. Finally, in order to quantitatively determine the dominant degeneracy for each individual galaxy in our sample we have defined the angle $\\theta_{i,j}$ like that satisfying \\begin{equation} \\cos \\theta_{i,j} = SIGN\\left({u_{i}\\over u_{j}}\\right) \\sqrt{u^{2}_{i}+u^{2}_{j}} \\end{equation} \\noindent where $u_{i}$ and $u_{j}$ are the $i$ and $j$ components of the PCA1 vector. The angle $\\theta_{i,j}$ simultaneously provides a measure of the angle between the PCA1 vector and the plane of physical properties $i,j$ and the sign of the degeneracy between the $i$ and $j$ properties. Thus, if $|\\cos\\theta_{i,j}|$$\\simeq$1 the degeneracy between the $i$ and $j$ properties would dominate the total degeneracy. Moreover, if $\\cos\\theta_{i,j}$$>$0 an increase in both the $i$ and $j$ properties would lead to similar observational properties, while if $\\cos\\theta_{i,j}$$<$0 the value of one of the properties should decrease. In Figure~\\ref{fig4}b we show the distribution of $\\cos\\theta_{i,j}$ as function of the age for the high-redshift sample assuming an error in the colors of 0.07\\,mag and the GALEX+SDSS+2MASS set available.  Along this section we will describe the results obtained from the analysis of the distributions shown in Figures~\\ref{fig3} \\& \\ref{fig4} for the different redshifts, observing errors, and band-pass combinations considered.  \\subsection{Nearby galaxies} \\label{sec4.1}  \\subsubsection{Formation Timescale} \\label{sec4.1.1}  With regard to the formation timescale in nearby galaxies, Figure~\\ref{fig5}a indicates that, even for relatively small observing errors, its uncertainty is very high (see also Table~\\ref{table3}) and shows a strong dependence with the value of the formation timescale itself. The larger uncertainty in the formation timescale for larger values of this quantity is mainly due to the small sensitivity of the optical-NIR colors of stellar populations with ages $t$$<<$$\\tau$ to changes in its formation timescale. The use of $U$-band data significantly reduces this uncertainty, probably due to the high sensitivity of this band to the presence of recent star formation that allows to rule out instantaneous-burst solutions when recent star formation associated with larger $\\tau$ values has effectively taken place. The use of NIR data, however, does not provide relevant information about the formation timescale of the stellar population. Moreover, the reduction achieved in the uncertainties of the different galaxy properties by using $JHK$ data compared with those obtained using exclusively $K$-band data is very small (see Table~\\ref{table3}). As we will show in Sections~\\ref{sec4.2} \\& \\ref{sec4.3} this is not the case for the intermediate and high-redshift galaxies, where these bands now cover the redshifted optical spectrum. Finally, in the same way that the $U$-band, the use of UV data provides an additional reduction in the formation timescale. As it is clearly seen in Figure~\\ref{fig2}, for a particular galaxy the formation timescale is mainly degenerate with the age of the stellar population, in the sense that, within the observing errors assumed, an increase in the formation timescale accompanied by an increase in the age can result in similar UV-optical-NIR colors. Due to this age-timescale degeneracy part of the reduction in the timescale uncertainty obtained by the use of UV data can be explained by the significant reduction in the age uncertainty achieved by including UV data (see below).  \\subsubsection{Age} \\label{sec4.1.2}  With respect to the age determination in nearby galaxies ($z$=0), the Figure~\\ref{fig5}a also shows that a significant reduction in the age uncertainty is achieved by including NIR data. It is important to note that the use of additional NIR data result in the same dominant degeneracy that if only optical data are used (see below), but the range of physical properties where this degeneracy takes place is significantly smaller. The most significant improvement in the age determination, however, is obtained when UV data are available. This is mainly due to the emission arising from post-AGB stars in low-metallicity populations and at ages younger than 10\\,Gyr and to the ``UV-upturn'' in high-metallicity evolved ($t$$>$10\\,Gyr) stellar populations that result in highly peculiar UV-optical colors. However, the uncertainty in the modeling of post-AGB stars (Charlot et a$.$ 1996) and the low-mass core helium-burning Horizontal Branch (HB hereafter) and evolved HB stars that lead to the ``UV-upturn'' (Yi et al$.$ 1997), introduce additional errors in the UV-optical colors during the data-models comparison. Charlot el al$.$ (1996) estimated using two different theoretical prescriptions that the uncertainty only in the post-AGB phase modeling could result in differences of about 1\\,mag in the UV-optical colors of a several-Gyr-old stellar population. Moreover, although we assumed the same observing errors for all the bands, the faint UV emission of evolved stellar populations is expected to result in very large observing errors in the UV-optical colors. Therefore, while the stellar evolution of these stars is not well understood the age determination in old stellar populations should not rely on the use of UV data.  Along with the formation timescale, the age of the stellar population in nearby galaxies is mainly degenerate with the dust extinction, in the sense that older stellar populations with low dust content have similar colors to highly-extincted, younger stellar populations. Although in the case of very old stellar populations the age-extinction degeneracy also competes with the age-metallicity degeneracy (see Worthey 1994), the age-extinction degeneracy is still dominant in this range for all the band-pass combinations and observing errors considered in this work. It could be argued that the strong discretization of the metallicity in our models could be responsible for the relatively weak age-metallicity degeneration derived. However, the fact that this behavior is observed even for the largest observing errors considered indicates that it is real and a natural consequence of the use of broad-band data. In this sense, the combination of broad-band with narrow- or medium-band data or spectroscopic indexes would break the age-extinction degeneracy, making of the age-metallicity the dominant degeneracy (see Worthey 1994).  \\subsubsection{Dust Extinction} \\label{sec4.1.3}  The dust extinction is derived with a very high accuracy ($E(B-V)$=0.04-0.20\\,mag) even for large observing errors and relatively low number of observables (see Table~\\ref{table3}). In the case of the nearby galaxies, the uncertainty in the dust extinction does not depend on the value of the dust extinction itself and is mainly degenerate with the age of the stellar population (see above) with some contribution from the extinction-metallicity degeneracy. In combination with $BVRI$ optical data either the use of UV, $U$, or NIR data provide a significant reduction in the dust-extinction uncertainty. In order to better derive the dust extinction the use of a wider wavelength baseline in wavelength (e.g$.$ using UVIJK) is more effective than fully covering the optical range (UBVRI). This is mainly due to the reduction in the metallicity uncertainty by the use of NIR data (see below) that leads, via the extinction-metallicity degeneracy, to a reduction in the dust-extinction uncertainty. Again, the use of $JHK$ data instead of only $K$-band data do not lead to a significant reduction in the dust-extinction uncertainties.  \\subsubsection{Metallicity} \\label{sec4.1.4}  With regard to the metallicity of the stellar population the uncertainties derived are strongly dependent on the band-pass combination available and the value of the metallicity itself. In particular, the uncertainties derived are smaller as the metallicity becomes higher (see Figure~\\ref{fig5}a). Within the age range considered, the main contributors to the optical and NIR emission of SSP galaxy are the main-sequence and RGB stars. However, for a more constant star formation, a significant contribution from core-Helium-burning stars is expected (see Charlot \\& Bruzual 1991). In order to determine the source of the metallicity dependence of these uncertainties we have produced the same diagrams shown in Figure~\\ref{fig5}a but restricted to formation timescales shorter than 50\\,Myr. The analysis of this diagram shows no dependence of the uncertainties with metallicity, which implies that the source of the dependence was the distinct photometric evolution of high-metallicity core Helium burning stars (Mowlavi et al$.$ 1998). It is worth noting, however, that at very high metallicities the uncertainties in the modeling of the stellar populations are themselves very large because of the lack of very metal rich stars of any age in the Solar neighborhood that could be used as spectral calibrators (see Charlot et al$.$ 1996).  The most significant reduction in the mean metallicity uncertainty is achieved when NIR data are used in combination with optical data (see Figure~\\ref{fig5}a). Although the uncertainties in the model predictions for the thermally pulsating AGB (TP-AGB hereafter) and the upper RGB can result in differences in the ($V-K$) color predicted by different models of $\\sim$0.10-0.15\\,mag (Charlot et al$.$ 1996), the improvement in the metallicity determination by the use of NIR data is still relevant. In this sense, in Table~\\ref{table3} we show that the mean metallicity uncertainty for the U+BVRI set is 0.32\\,dex assuming an observing error of 0.03\\,mag, while the uncertainty for the U+BVRI+K set assuming an observing error of 0.10\\,mag is only 0.26\\,dex.  \\subsubsection{Stellar Mass} \\label{sec4.1.5}  As input for the $K$-band mass-to-light ratio of the stellar populations we have adopted M$_{K,\\odot}$=3.33 (Worthey 1994). It should be noticed that along with the errors in the stellar mass-to-light ratios derived here the misunderstanding of the actual IMF introduce an additional, systematic uncertainty, which, in fact, constitutes the most important source of error in the determination of the galaxy stellar mass (Bell \\& de Jong 2001). In addition, the poor constraints on the theoretical isochrones of upper-RGB stars and AGB stars can result in a 20~per cent uncertainty in the $K$-band mass-to-light ratio (Charlot et al$.$ 1996). In Figures~\\ref{fig6}a \\& \\ref{fig6}b we show the uncertainties expected in the $K$-band mass-to-light ratio from different sets of observables that include $K$-band data. These uncertainties show a strong dependence with the galaxy age and formation timescale in the sense that larger uncertainties are expected at lower values of the formation timescale and older ages. Figure~\\ref{fig5}a shows that the value of the age uncertainty (in log~$t$ scale) is almost independent of the age itself. In addition, Figure~\\ref{fig1} indicates that the rate of change in the $K$-band mass-to-light ratio (with log~$t$) is higher when the stellar population becomes older, specially for very low values of the formation timescale. Therefore, for a constant uncertainty in log~$t$, an increase in the uncertainty of the mass-to-light ratio at very old ages is expected.  Figure~\\ref{fig6}a also shows that the mass-to-light ratio determination is biased toward lower values. This bias, which is particularly important at old ages, is probably due to the upper limit of 15\\,Gyr in age imposed during the data-models comparison (see Table~\\ref{table1}), although other contributors can not be ruled out (see Sections~\\ref{sec4.2.2} \\& \\ref{sec4.3.2}). As it is clearly seen in Figure~\\ref{fig6}b, the use of UV data allows to reduce both the uncertainty and bias in the mass-to-light determination. This reduction is directly related with the reduction in the age uncertainty described above. However, as we already commented, the use of UV data for the study of stellar populations with ages older than several Gyr can lead to wrong conclusions because of the uncertain modeling of the post-AGB phase and the ``UV-upturn''.  The behavior described above for the timescale, age, dust extinction, metallicity, and stellar mass is identical for any observing error but with larger mean uncertainties for larger observing errors. The reader is referred to the Table~\\ref{table3} for the dependence of the mean uncertainties in the different galaxy properties derived with the observing errors.  \\subsection{Intermediate-redshift galaxies} \\label{sec4.2}  \\subsubsection{Formation Timescale} \\label{sec4.2.1}  In Figure~\\ref{fig7} we show the uncertainties derived for the properties of intermediate-redshift galaxies ($z$=0.7). With regard to the formation timescale the uncertainties are very large (2-3\\,Gyr), even larger than for the nearby galaxy sample. As we commented in Section~\\ref{sec4.1}, the optical-NIR colors are quite insensitive to changes in the formation timescale with $t$$<<$$\\tau$. Therefore, since we are assuming that these galaxies are statistically younger than the those observed in our Local Universe (see Table~\\ref{table1}) and the range in formation timescale is obviously the same, the uncertainty in the formation timescale is necessarily higher. For the same reason the uncertainty at very low timescale values is much lower than at high timescale values.  The upper panel of Figure~\\ref{fig7}a also suggests a significant bias in the timescale determination toward lower values of this property. This bias is also the consequence of the small changes in the optical-NIR colors of these galaxies with the timescale when the age is younger than the timescale value. In this case, the higher rate of change in the colors toward lower formation timescales systematically leads to lower values in order to reproduce the probability distribution associated with the observing errors. It is worth noting that, because of the reduction of this bias, the use of a larger number of bands may result in some cases in a higher timescale uncertainty (see Table~\\ref{table3} for the results on the U+BVRI+K and U+BVRI+JHK sets). Like in the nearby galaxies case, the dominant degeneracy involving the formation timescale is the age-timescale degeneracy, in the sense that older galaxies with high formation timescales have similar colors that younger galaxies with a more instantaneous star formation. This is true for any band-pass combination considered. With regard to the optimal set of observables, Table~\\ref{table3} demonstrates that for the same number of bands the use of wider wavelength baselines results in lower uncertainties. In particular, the use of the UVIJK set reduces the timescale, age, and metallicity uncertainties inherent to the UBVRI set providing also a much lower dust-extinction uncertainty than the BVRI+K set. On the other hand, the SDSS+2MASS and GALEX+SDSS+2MASS sets result in very similar uncertainties (see Table~\\ref{table3}), which implies that the optical and NIR bands provide most of the information available in the UV and in the blue part of the optical spectrum about the galaxy age, star formation history, and metallicity.  \\subsubsection{Age} \\label{sec4.2.2}  With respect to age of the intermediate-redshift galaxies the uncertainties derived are significantly smaller than in the nearby-galaxies case. This is mainly due to the higher rate of change in the rest-frame optical colors within the age range assumed for these galaxies compared with that assumed for the nearby galaxies (see Table~\\ref{table1} and Figure~\\ref{fig1}). In addition, the fact that the $K$-band now corresponds to the rest-frame $J$-band emission implies that the effect of the uncertainties in the model predictions associated with the upper RGB and AGB evolutionary stages is less important (see Section~\\ref{sec4.4.2}). On the other hand, the use of $U$-band data for determining ages older than 1\\,Gyr at these redshifts is strongly limited by the uncertainty in the modeling of the rest-frame UV emission from post-AGB stars (Charlot et al$.$ 1996; see Section~\\ref{sec4.1}). However, the most significant decrease in the age uncertainty is achieved when NIR data are used, specially if data in all the bands ($JHK$) are available. This is probably due to the fact that the $JHK$ set provides information simultaneously about the presence of AGB stars (via the rest-frame $z'$ and $J$ bands) and main-sequence stars (via the rest-frame $R$-band).  Figure~\\ref{fig5}a also shows the existence of a significant bias toward younger ages for the BVRI and UBVRI sets. In this case the presence of this bias is due (1) to the existence of a formation timescale bias and a strong age-timescale degeneracy and (2) to the fact that the optical colors of the stellar populations change more slowly as the population becomes older. In the latter case, in order to reproduce the distribution of optical colors associated with the observing errors, the best-fitting solution should be found at younger ages, where the intrinsic dispersion of the model colors is larger. As we show below a bias in age also results in a bias in the galaxy $K$-band mass-to-light ratio. The use of wider wavelength baselines allows to significantly reduce this bias. In particular, the use of the UVIJK leads to a less severe bias and lower age uncertainties than the U+BVRI and the U+BVRI+K sets. Within the age uncertainty interval the degeneracy is mainly dominated by the age-timescale degeneracy with some contribution from the age-extinction degeneracy in those band-pass combinations that do not include UV or $U$-band data.  \\subsubsection{Dust Extinction} \\label{sec4.2.3}  The dust extinction in the sample of intermediate-redshift galaxies is derived with high accuracy, specially when $U$-band data are available (see Figure~\\ref{fig7}a). In this case, the uncertainties expected in $E(B-V)$ are in any case smaller than 0.10\\,mag for observing uncertainties as high as $\\Delta C_n$=0.10\\,mag. The significant reduction achieved, if we compare these results with those derived for the nearby galaxies, is due to the very high sensitivity of the redshifted UV emission to the presence of small amounts of dust. In those band-pass combinations not including $U$-band data we notice a clear dependence of the dust-extinction uncertainty with the value of the extinction itself, with larger uncertainties at larger values of the extinction (see Figure~\\ref{fig7}a). The analysis of the PCA1 components also indicates that at dust-extinction values higher than $E(B-V)$$>$0.5\\,mag the age-extinction degeneracy becomes very important. This implies that in highly extincted intermediate-redshift galaxies a small increase in the amount of dust can lead to the same optical colors (specially if $U$-band data are not used) that a comparable decrease in the age of the stellar population would produce.  \\subsubsection{Metallicity} \\label{sec4.2.4}  With regard to the metallicity uncertainty, Figure~\\ref{fig7}a shows that the uncertainty decreases with the value of the metallicity itself. The reduction is particularly important when NIR data are available. The use of the three $JHK$ NIR bands reduces this uncertainty over the whole range of metallicities. In this sense, the use of the UVIJK set results in lower metallicity uncertainties than the U+BVRI and the U+BVRI+K sets (see Table~\\ref{table3}). It is important to keep in mind that the $JHK$ filters now cover the rest-frame $R$, $z'$, and $J$ bands. In the age range considered the main contribution to the rest-frame optical emission comes from main-sequence stars. On the other hand, the rest-frame NIR emission, along with main-sequence stars, shows an important contribution from AGB and core-Helium-burning stars (see Charlot \\& Bruzual 1991). The role played by AGB stars is more relevant if the formation is instantaneous, while the core-Helium-burning stars may dominate the total NIR emission for a more constant star formation scenario. Therefore, the behavior described above is probably due to the distinct evolution of high-metallicity AGB stars (see Willson 2000 and references therein) and core-Helium-burning stars (Mowlavi et al$.$ 1998) compared with the relatively well-defined sequence in their evolutionary properties established for sub-solar metallicities. Within the error intervals derived, the metallicity is mainly degenerate with the age, specially in those sets including $U$-band data. This is probably due to the reduction in the age-extinction degeneracy thanks to the information provided by the $U$-band data about the rest-frame UV.  \\subsubsection{Stellar Mass} \\label{sec4.2.5}  The comparison between Figures~\\ref{fig6}b and \\ref{fig7}b shows that the mean uncertainties in the $K$-band mass-to-light ratio (or stellar mass) of intermediate-redshift galaxies are much lower than those derived for the nearby sample. First, it is important to note that in these figures we represent absolute errors. For a Solar-abundant 12\\,Gyr-old nearby galaxy formed instantaneously the $K$-band mass-to-light ratio is $\\sim$1.3\\,M$_{\\odot}$/L$_{K,\\odot}$, while for a 5\\,Gyr-old galaxy at $z$=0.7 is $\\sim$0.4\\,M$_{\\odot}$/L$_{K,\\odot}$. Therefore, the relative uncertainties, assuming the average absolute uncertainties given in Table~\\ref{table3} for $\\Delta C_n$=0.07\\,mag, would be about 30 and 20~per cent, respectively for the nearby and intermediate-redshift galaxies. Although this still implies a significant improvement in the $K$-band mass-to-light ratio determination, it is also noticeable that the $K$ filter now traces the rest-frame $J$-band luminosity, which is more affected by the misunderstanding about the actual IMF (see Bell \\& de Jong 2001). Finally, the $J$-band luminosity is also more sensitive to small differences between the assumed exponential star formation and the galaxy actual star formation history than the rest-frame $K$-band data.  As we pointed out in Section~\\ref{sec4.2.2} the bias in the age determination toward lower age values also leads to a strong bias in the $K$-band mass-to-light ratio of intermediate-redshift galaxies due to the systemic decrease in the rest-frame $J$-band luminosity per Solar mass with the age of stellar population when the age is older than $\\sim$1\\,Gyr.  \\subsection{High-redshift galaxies} \\label{sec4.3}  \\subsubsection{Formation Timescale} \\label{sec4.3.1}  With regard to the formation timescale, Figure~\\ref{fig8}a shows that the bias toward lower timescale values observed at intermediate redshift is even more pronounced at high-redshift. This bias is a natural consequence of the difficulty of deriving/predicting the long-term star formation history of a galaxy when it is still very young. This is also evidenced by the fact that the mean timescale uncertainty increases systematically with redshift for the same observing errors and band-pass combinations. In Table~\\ref{table3} we also show that in many cases (BVRI vs$.$ U+BVRI; SDSS vs$.$ SDSS+2MASS) the mean timescale uncertainties increase when a larger number of observing bands is used, with a progressive reduction in this bias. As in the intermediate-redshift case the dominant degeneracy involving the galaxy formation timescale occurs with the age of the stellar population.  \\subsubsection{Age} \\label{sec4.3.2}  The large formation timescale uncertainty described above and the existence of a strong age-timescale degeneracy, specially at ages older than 100\\,Myr, lead to very large age uncertainties, even larger than those derived for the intermediate-redshift galaxies.  The age-timescale degeneracy at ages younger than 100\\,Myr is significantly smaller because at these young ages a change in the formation timescale, which ranges between 200\\,Myr-6\\,Gyr (see Table~\\ref{table1}), does not affect to the UV-optical-NIR colors of the stellar population. In other words, the degeneracy in timescale within this age range is complete and no correlation between the age uncertainty and any other uncertainty is expected. In this case the main degeneracies are the age-extinction and the age-metallicity ones.  Moreover, the age-timescale degeneracy in combination with the bias in formation timescale described above are also responsible for the strong bias in age observed in Figure~\\ref{fig8}a at ages older than $\\sim$50\\,Myr. The fact that the UVIJK set provides a better age and timescale determination than the U+BVRI and U+BVRI+K sets demonstrates the importance of obtaining $JHK$ data in order to derive the properties of high-redshift galaxies. This is due to the fact that the $JHK$ filters now cover the rest-frame $V$, $R$, and $z'$ optical bands, where the changes due to galaxy evolution are more noteworthy and the information content about the galaxy properties is larger. In particular, the $JHK$ filters would provide information simultaneously about the presence of main-sequence stars (via the rest-frame $V$ and $R$ bands), core-Helium-burning stars (via the rest-frame $V$, $R$, and $z'$ bands), and AGB stars (via the rest-frame $z'$-band; $t$$>$0.5\\,Gyr).  \\subsubsection{Dust Extinction} \\label{sec4.3.3}  Because of the extensive coverage of the UV range of the spectrum, the study of high-redshift galaxies using optical-NIR colors leads to very small dust-extinction uncertainties. In this sense, the dust-extinction uncertainties given in Table~\\ref{table3} at this redshift assuming an observing error of 0.10\\,mag are in the range $E(B-V)$=0.03-0.07\\,mag. The dust extinction within the interval of uncertainty is mainly degenerate with the age of the stellar population.  \\subsubsection{Metallicity} \\label{sec4.3.4}  Figure~\\ref{fig8} shows that the metallicity uncertainty for the high-redshift sample does not show the strong metallicity dependence found in nearby and intermediate-redshift samples. Only when $JHK$ NIR data are available the uncertainties at very high metallicities become significantly smaller than those derived for the low metallicity galaxies. As we commented in Section~\\ref{sec4.2.4} for the intermediate-redshift case, this is probably due to the distinct signature of high-metallicity core-Helium-burning stars (eg$.$ in the number ratio of blue-to-red supergiants; Mowlavi et al$.$ 1998) within the age range considered. In the case of a SSP galaxy, these stars dominate the rest-frame $VRz'$ ($JHK$ at $z$=1.4) emission for ages younger than 0.4\\,Gyr, while the emission at shorter wavelengths comes mainly from main-sequence stars (see Charlot \\& Bruzual 1991). It is important to note that the core-Helium-burning stars may dominate the emission in the $R$ and $z'$ bands up to ages of 5\\,Gyr for larger formation timescales. Within the uncertainty intervals obtained, the metallicity is mainly degenerate with the age of the stellar population.  \\subsubsection{Stellar Mass} \\label{sec4.3.5}  The $K$-band mass-to-light ratio uncertainties derived here are very small compared with those obtained from the nearby and intermediate-redshift samples, with values ranging between 0.01 and 0.06\\,M$_{\\odot}$/L$_{K,\\odot}$. If we adopt a $K$-band mass-to-light of 0.27\\,M$_{\\odot}$/L$_{K,\\odot}$, which corresponds to the value expected for a 3\\,Gyr-old galaxy with Solar metallicity, the relative uncertainty would range between 5\\% and 20\\%, depending of the band-pass combination available. Figure~\\ref{fig8}b shows that there is also a strong dependence of the mass-to-light ratio uncertainty with the value of the mass-to-light ratio itself. In particular, a clear minimum in its uncertainty is observed at ages older than 8\\,Myr, which is probably associated with the evolution of the massive stars off the main sequence toward the red supergiant phase. During this part of the evolution a sudden change in the rest-frame $z'$ luminosity and optical colors of a SSP is produced, which could explain why the uncertainty is particularly small around this age value.  \\subsection{Effects of the Model Uncertainties} \\label{sec4.4}  In this section we analyze the results obtained when the optical-NIR colors of a sample of galaxies generated using the GISSEL99 models are compared with the predictions of the P\\'{E}GASE evolutionary synthesis models. We have restricted this comparison to the nearby sample and the range of properties specified in Table~\\ref{table1}. The results of this comparison are shown in Figure~\\ref{fig9}.  \\subsubsection{Formation Timescale} \\label{sec4.4.1}  Figure~\\ref{fig9}a shows that the same bias toward lower values of the formation timescale that we noted for the intermediate and high-redshift samples is also present in this case (see Sections~\\ref{sec4.2.1} \\& \\ref{sec4.3.1}). The main reason for the existence of this bias is the small change in the optical-NIR colors of the stellar population with the timescale when the age $t$$<<$$\\tau$. Therefore, in order to compensate both the observing errors and the differences in the color predictions between the GISSEL99 and P\\'{E}GASE models, the best-fitting solution has to be found at lower values of the timescale where the intrinsic dispersion of the colors is larger. The existence of this strong bias also leads to very small timescale uncertainties compared with those obtained using the GISSEL99 models. Within the uncertainty intervals derived, the dominant degeneracy involving the galaxy formation timescale is the age-timescale degeneracy.   \\subsubsection{Age} \\label{sec4.4.2}  With regard to the age determination, the uncertainties derived are very similar for the BVRI and U+BVRI sets. However, for those band-pass combinations including NIR data the ages derived are strongly biased toward younger ages. The reason for this bias, which also leads to significantly smaller age uncertainties, is the offset in the ($J-H$) and ($H-K$) colors between the GISSEL99 and the P\\'{E}GASE model predictions (see Figure~\\ref{fig1}) due to the differences in the modelling of the upper RGB and AGB phases. In particular, Figure~\\ref{fig1}b shows that the P\\'{E}GASE models are $\\sim$0.07\\,mag redder in ($J-H$) and $\\sim$0.04\\,mag redder in ($H-K$) than the GISSEL99 models within the age range 4-12\\,Gyr. Therefore, in order to compensate for this difference in color, the best-fitting solution usually leads to younger ages, which within this age range imply bluer colors. Because the differences in the colors between the two models only occur in the NIR, the optical colors predicted by the P\\'{E}GASE models at these younger ages should be bluer than those of the sample. Therefore, in order to compensate for this effect, the age bias described above has to be accompanied by strong biases in dust extinction and/or metallicity that would lead to redder optical colors. Within the error intervals derived the total degeneracy is dominated by the age-timescale and age-extinction degeneracies.  \\subsubsection{Dust Extinction} \\label{sec4.4.3}  As we commented above (see also Figure~\\ref{fig9}a) there is a strong bias in dust extinction estimates toward higher extinction values when NIR data are used. This bias, along with the metallicity bias described in Section~\\ref{sec4.4.4}, results in a global reddening of the optical colors but a small change in the NIR colors of the galaxies in the sample. On the other hand, at very high extinction values the uncertainties are also biased by the upper limit in $E(B-V)$ imposed during the data-models comparison procedure (see Table~\\ref{table1}). The mean uncertainties derived, both in age, dust extinction, and metallicity are very similar to those obtained by using the GISSEL99 models.  \\subsubsection{Metallicity} \\label{sec4.4.4}  The distribution of the uncertainty in metallicity shown in Figure~\\ref{fig9}a indicates that a strong bias toward higher metallicity values is present when NIR data are available. As we commented in Section~\\ref{sec4.4.2}, this bias is probably related with the age bias and the differences in the NIR colors predicted by the two sets of models. As in the case of GISELL~99 models, the comparison with the P\\'{E}GASE models leads to a clear dependence of the uncertainty with the metallicity value itself, with smaller uncertainties at very high metallicities (see Section~\\ref{sec4.1.4}).  \\subsubsection{Stellar Mass} \\label{sec4.4.5}  The results shown in Figure~\\ref{fig9}b with regard to the $K$-band mass-to-light ratios mainly reflect the biases in the galaxy property determination, with the stellar masses derived systematically smaller than the input values. This is due (1) to the bias toward younger ages described in Section~\\ref{sec4.4.2} and (2) to the higher $K$-band luminosity per unit mass of the P\\'{E}GASE models compared with the GISSEL99 models (see Figure~\\ref{fig1}b). Because of the stronger bias in age, the mean uncertaintines in the $K$-band mass-to-light ratio are smaller than those obtained using the GISSEL99 models (see Section~\\ref{sec4.1.5}). Finally, Figure~\\ref{fig9}b shows that the mass-to-light ratio uncertainty becomes higher at older ages and lower timescale values. This behavior, which is also present in the case of the GISSEL99 models (see Section~\\ref{sec4.1.5}), is due to the progressive increase in the rate of change of the $K$-band mass-to-light ratio with log~$t$ (see Figure~\\ref{fig1}) accompanied by a small dependence of the age uncertainty (in log~$t$ scale) with the value of the age itself. ",
        "conclusions": "\\label{sec5}  In this study we have analyzed the dependence of the uncertainties and degeneracies in the galaxy properties upon different parameters: (1) the combination of bands available, (2) the observing errors, and (3) the galaxy properties themselves (including redshift).  Here we summarize our main results and point out some directions for the optimization of galaxy evolution studies using broad-band photometry data. We describe separately the nearby, intermediate, and high-redshift cases.  {\\bf Nearby galaxies:} In order to determine the star formation history, age, and dust extinction of nearby galaxies with relatively small uncertainties the use of $U$-band data is fundamental. The availability of $K$-band data also allows a reduction in the uncertainty in the age and metallicity of the stellar population, but the use of additional $J$ and $H$-band data is largely redundant. The use of the $K$-band data is unfortunately limited by the existence of large uncertainties in the modeling of the $K$-band luminosities and NIR colors of stellar populations. The most significant reduction in the age and $K$-band mass-to-light ratio uncertainty is achieved when UV data are used. The poor treatment of the post-AGB and ``extreme'' HB phases by the existing evolutionary synthesis models introduce, however, an additional uncertainty during the data-model comparison, which is particularly important in the case of very old stellar populations. For the same number of observing bands, the availability of wider wavelength baselines results in lower uncertainties. Both the formation timescale and $K$-band mass-to-light ratio uncertainties are larger when the corresponding values for these properties are larger. On the other hand, the metallicity uncertainty decreases with the value of the metallicity itself due to the distinct photometric evolution of high-metallicity core Helium burning stars.  A complete description of the physical reasons behind these conclusions and of the degeneracies responsible for the uncertainties described above are given in Section~\\ref{sec4.1} (see also Section~\\ref{sec4.4}).  {\\bf Intermediate-redshift galaxies:} The star formation history of intermediate-redshift galaxies can be derived with worse precision than in nearby galaxies because their stellar populations are younger. The age uncertainty is smaller than in the nearby-galaxies case and shows a strong bias toward younger ages. A significant reduction of this bias and of the mean uncertainties is achieved when NIR data are used, especially if all three $J$, $H$, and $K$-band data are available. The dust-extinction uncertainty is larger for larger values of the dust extinction itself. The use of $U$-band data provides an important reduction of this dependence and of the mean dust-extinction uncertainty. If $U$-band data are available the use of additional UV data do not provide much more information about the galaxy properties. The use of NIR data ($J$, $H$, and $K$-band data) significantly reduces the uncertainty in the metallicity of the galaxy. The absolute and relative uncertainties in the galaxy $K$-band mass-to-light ratio are smaller than those derived for nearby galaxies. However, the fact that the $K$ filter now covers the rest-frame $J$-band leads to a larger uncertainty associated with the IMF and with the parameterization of the galaxy star formation history and, consequently, to a larger stellar mass uncertainty. For a more detailed description see Section~\\ref{sec4.2}.  {\\bf High-redshift galaxies:} As expected, the bias and mean uncertainty in the determination of the timescale for the galaxy formation are even larger in this case that in the nearby or intermediate-redshift galaxies. The age of the stellar population is derived with a large uncertainty, only reduced when $JHK$ data are available. The dust-extinction in these galaxies can be derived to a very high accuracy even when only optical data are available. The use of $JHK$ data is fundamental in order to improve both the age and metallicity determinations. A complete description of the uncertainties and degeneracies between these properties is given in Section~\\ref{sec4.3}.  Some of the conclusions drawn above can also be found through the literature expressed in a qualitative way. However, this work constitutes the first systematic and quantitative study on the optimization of broad-band photometry for studies on the evolution of galaxies. It is important to note that the application of these results to future galaxy surveys can help to reduce the uncertainty in the derivation of the galaxy physical properties, sometimes by weakening a particular degeneracy but most of the time by decreasing the intervals over which this degeneracy takes place. In Table~\\ref{table3} we have summarized the mean uncertainties in the galaxy stellar population properties derived in this paper considering different redshifts, sets of observables, and observing errors.  Our results are directly applicable to spectrophotometric surveys like the SDSS and surveys looking for emission-line galaxies at fixed redshifts (Martin, Lotz \\& Ferguson 2000; Moorwood, van der Werf, Cuby \\& Oliva 2000; Iwamuro et al$.$ 2000; Pascual et al$.$ 2001; Zamorano et al$.$, in preparation). However, in the case of the blind-redshift surveys a comparison between our results and those from previous studies on the optimization of the photometric-redshifts technique (Kodama et al$.$ 1999; Bolzonella et al$.$ 2000; Mobasher \\& Mazzei 2000; Wolf et al$.$ 2001) is still needed."
    },
    "0201/astro-ph0201296_arXiv.txt": {
        "abstract": "Dodson \\& Ellingsen (2002) included several observations with significant pointing errors, invalidating the upper limits found in these directions. These have now been reobserved or recalculated. A new table of upper limits has been generated, and two more masers that would have been seen have been found. ",
        "introduction": "Our followup to the original observations \\cite{DE02}, performed on 15 December 2003, with the ATCA in configuration 6A, had only 10 baselines so the detection limits are a little higher than the original data set and are listed in Table 1. As before the detection peak heights, velocities and widths are found from a Gaussian fit to the unsmoothed data. The detections are listed in Table 2.   \\subsubsection*{G339.622-0.121}    The 4765-MHz maser emission from this new detection peaks at -37.0~\\kms . This site has 1720-, 1665- and 6035-MHz OH masers \\cite{C99}. There is a nearby 1612-MHz emission site which also displays methanol masing. In this region there is no detectable 4765-MHz emission. The flux from this source is below the limit of Smits \\shortcite{S03}.  \\subsubsection*{G347.628+0.148}   This source was discovered by Smits \\shortcite{S03}, with a slightly stronger flux density (a peak of 5 Jy). As the source is very narrow (0.3 km/s FWHM) our observations may underestimate the strength, but the measured emission was falling during the Hartsebethock observations and is consistent. This source has a peak velocity of -96.3 \\kms\\ and a secondary peak of 0.3~Jy at -95.1 \\kms\\ from the same position. This source also has emission from 1612-, 1665-, 1720- and 6035-MHz, and 6.7-GHz methanol. It is an \\UCHII\\ region.   \\begin{table*} \\begin{minipage}{\\textwidth} \\caption{Sources searched for 4765-MHz OH maser emission. The 5-$\\sigma$ limit column is for emission in a vector averaged spectrum at centred on the position, and covering the velocity range listed. References : a = Caswell \\protect\\shortcite{C97} ; c = Caswell \\protect\\shortcite{C99} ; d = Cohen \\etal\\/ \\protect\\shortcite{CMC95} ; } \\begin{tabular}{lcccrcl} \\hline & \\multicolumn{2}{c}{\\bf Reference position} & {\\bf Velocity} & {\\bf 5-$\\sigma$} & {\\bf Central}  & \\\\ {\\bf Source} & {\\bf Right Ascension} & {\\bf Declination} & {\\bf Range}    & {\\bf Limit}           & {\\bf Velocity} & \\\\ {\\bf Name}   & {\\bf (J2000)}         & {\\bf (J2000)}     & {\\bf (\\kms)}   & {\\bf (Jy)}            & {\\bf (\\kms)}   & {\\bf References} \\\\ \\hline \\multicolumn{7}{c}{reobserved}\\\\ G339.622$-$0.121& 16:46:06.030 & -45:36:43.90& -130--95 & 0.11 & -36.7    & c  \\\\%1720 G347.628$+$0.149& 17:11:50.888 & -39:09:29.00& -130--95 & 0.12 & -96.7     & a  \\\\ G348.550$-$0.979& 17:19:20.418 & -39:03:51.65& -130--95 & 0.12 & -13.0,-12.1&a,c\\\\%1720 G351.581$-$0.353& 17:25:25.085 & -36:12:46.08& -130--95 & 0.15 & -93.8     & a  \\\\ \\multicolumn{7}{c}{corrected}\\\\ G351.417$+$0.645& 17:20:53.370 & -35:47:01.16& -106--135 & 0.11 & -10.0,-10.4&a,c\\\\%1720 {\\em IRAS}17175$-$3544& 17:20:54.945 & -35:47:02.38& -106--135 & 0.11 & -10.0     & d  \\\\ G351.775$-$0.536& 17:26:42.559 & -36:09:15.99& -106--135 & 0.15 & -8.7,4    & a,c  \\\\%1720 \\hline \\end{tabular} \\label{tab:nondet4765} \\end{minipage} \\end{table*}  \\begin{table*} \\begin{minipage}{\\textwidth} \\caption{Sources with detected 4765-MHz OH emission. References : *=new source, e=Smits \\protect\\shortcite{S03};} \\begin{tabular}{lcccccrl} \\hline & \\multicolumn{2}{c}{\\bf 4765-MHz Maser}    & {\\bf Peak Flux} & {\\bf Velocity} & {\\bf Width } &                   \\\\ {\\bf Source} & {\\bf Right Ascension} & {\\bf Declination} & {\\bf Density}   & {\\bf of Peak}  & {\\bf of Peak}    &                   \\\\ {\\bf Name}   & {\\bf (J2000)}         & {\\bf (J2000)}     & {\\bf (Jy)}      & {\\bf (\\kms)}   & {\\bf (\\kms)}   &  {\\bf References} \\\\ \\hline G339.622$-$0.121&16:46:06.03$\\pm0.04$ & -45:36:43.3$\\pm0.2$ & 0.56 & -37.0 & 0.5 &   * \\\\ G347.628$+$0.149&17:11:51.01$\\pm0.02$ & -39:09:29.0$\\pm0.5$ & 3.65 & -96.3 & 0.3 &   e \\\\ \\hline \\end{tabular} \\label{tab:det4765} \\end{minipage} \\end{table*}  \\begin{figure*} \\dfig{g339.622-0.121}{g347.628+00.149} \\caption{Total intensity spectra of the 4765~MHz OH masers reobserved with the ATCA} \\end{figure*} ",
        "conclusions": "Both sources found also display 1720- and 6035-MHz OH emission, so they do not significantly alter our previous conclusions about the association of 4765-MHz OH emission with these two transitions."
    },
    "0201/astro-ph0201543_arXiv.txt": {
        "abstract": "Models for the origin of the hard X-ray background have suggested that sources with the most accretion activity lie hidden in highly obscured AGN. We report on our study of hard, serendipitous sources in the fields of \\c\\ clusters with fluxes close to the turn-over in the source-counts. These include two Type II quasars with measured X-ray luminosities $>10^{45}$ \\ergps\\ and column-densities $> 10^{23}$ cm$^{-2}$, one possibly being Compton-thick. Both show indications of redshifted Fe K$\\alpha$ line emission. Radiative transfer modelling of the broad-band spectrum of a highly-magnified source with deep {\\sl ISOCAM} detections implies the presence of warm-to-hot dust obscuring a quasar with L$_{\\rm UV}>10^{45}$~\\ergps. Multi-wavelength spectroscopic and photometric follow-up of the optically-faint sources suggests that these objects are found in the centres of massive evolved galaxies at a range of redshifts, with red optical / near-infrared colours dominated by the host galaxy. Detailed source identification is difficult due to the paucity of strong emission features, especially in the near-infrared. We present the main results from a sample of near-infrared spectra of optically-faint sources obtained with 4~m and 8~m telescopes. Through the study of the harder and brighter X-ray background source population, we are likely to be viewing the most intense phase of the growth of supermassive black holes. ",
        "introduction": "The \\c\\ observatory has largely resolved the hard (2--10 keV) X-ray background (HXRB) within two years of its launch, after almost four decades of scientific effort (\\cite{giacconi}; \\cite{mushotzky}). Follow-up work within the past year has revealed that this \\c\\ source population can be broadly split into quasars (a small fraction), narrow emission-line AGN, optically-normal galaxies with no sign of activity other than in X-rays and optically-faint sources which are difficult to identify (\\cite{bargeretal}; \\cite{giacconi2001} \\cite{alexanderfaint}; \\cite{willott}). {\\sl XMM-Newton}, with its higher effective area at 10 keV, has also begun to deliver results (\\cite{hasinger}; \\cite{barconsAXIS}) and the essential \\c\\ findings are confirmed with many new Type II AGN being found (eg, \\cite{lehmannxmm}).  The individual X-ray spectra of the hard sources are flat enough to account for the HXRB spectral slope of 0.4 (\\cite{marshall}; \\cite{gruber}), and the integrated flux can contribute between $\\sim 70-105$\\% of the HXRB intensity. The ambiguity in the absolute known intensity is still to be resolved. Its cause is cross-calibration mismatches between various X-ray missions and/or cosmic variance of the sources themselves (eg, \\cite{barcons}; \\cite{cowie_number_counts}). Essentially, these observations bear out the main prediction of models which synthesize the HXRB through the integrated emission of a population of obscured active galaxies (AGN) with a distribution of absorbing columns and spread over redshift (\\cite{setti}; \\cite{comastri}, \\cite{wfn}). The hard spectral slope is produced by the presence of large obscuring columns which erode the soft photons. This also has the exciting corollary that \\lq Type II\\rq\\ QSOs -- the missing population of powerful, absorbed sources should emerge in large numbers.  In fact, the optically-faint sources are prime candidates for being members of this Type II population. Due to their faintness, it has been difficult to determine their nature even with 10 m class telescopes (eg, \\cite{bargeretal}). Speculations include the possibility that these are very high redshift ($>$6) quasars, though \\cite*{alexanderfaint} find that such objects can be at most a small fraction of the total population.  We have been studying the optically-faint population of hard sources in \\c\\ cluster fields. The main aim-point of the observation is the cluster itself and typical exposure times range from 10--40~ks. We thus have $0.5-7$~keV flux coverage for point background sources of $3\\times 10^{-15}\\sim 10^{-13}$~\\ergpspsqcm. The contribution to the XRB is maximized at a turn-over in the source counts at $\\approx 10^{-14}$~\\ergpspsqcm\\ and we thus study the dominant flux contributors to the XRB by targeting this regime (eg, \\cite{cowie_number_counts}). It is likely that $40\\sim 70\\%$ of the 2--10~keV XRB is emitted from sources at such fluxes.  Unlike very deep, small-area surveys, our minimal selection criterion is absence or faintness on the Digitized Sky Survey\\footnote{http://archive.stsci.edu/dss/}. In this way we are able to select hard X-ray sources with weak optical counterparts from almost anywhere in the sky. We find that many of these sources are relatively bright and readily detected in the near-infrared (NIR), with \\k$\\sim 17$ (\\cite{c01a}, \\cite{moriond}). The large optical--IR colours (reaching as much as  $R-K\\approx 8$; see also \\cite{alexanderred}) are consistent with the flux being dominated by light from the host galaxy and any central AGN being highly obscured. Photometric redshift (\\zp; \\cite{hyperz}) estimates suggest that these are early-type galaxies at redshifts $z\\sim 1-2$, with a tail of sources at higher redshifts. One such highly obscured source was confirmed to be a Type~II QSO in the field of Abell~2390 (\\cite{c01b} and also discussed in section 2 of this work).  Herein we report the discovery of another Type II QSO in the field of the Abell~963 cluster. We also find potential redshifted Fe K$\\alpha$ lines in both the above Type~II QSOs. In addition, we report on our findings of near-infrared spectra of hard serendipitous \\c\\ sources. Through 4~m and 8~m telescope observations, we have found that few sources contain strong detectable emission lines and fewer still can be classified with confidence. This is most likely due to depletion / destruction of emission lines by large column-density gas and the associated dust. We also report on indications of clustering of hard \\c\\ sources in the field of A\\,963.  Quoted cosmological quantities assume H$_0=50$ km~s$^{-1}$ Mpc$^{-1}$ and q$_0=0.5$ throughout. ",
        "conclusions": "Through a dedicated campaign of targeting the hardest and brightest sources in \\c\\ cluster fields, we have been able to detect 2 Type II QSOs (with L$_{2-10}>10^{45}$~\\ergps\\ and \\nh$>10^{23}$~cm$^{-2}$. In one case, this column is probably as high as $10^{24}$~cm$^{-2}$. Both these sources show indications of strong Fe K$\\alpha$ lines, though we cannot constrain the line parameters with the current data. Through radiative-transfer modelling based on {\\sl ISOCAM} detections of another source, we infer the presence of warm-to-hot dust obscuring a central AGN with L$_{\\rm UV}>10^{45}$~\\ergps, being reprocessed into the mid-to-far infrared. We detect at least one source with an anomalous dust:gas ratio. Most sources are consistent with Compton-thin Seyfert IIs, with the redshift distribution (both spectroscopic measurements and photometric estimates) peaking at $z\\approx 1$ and a significant tail to higher redshifts. These sources are optically-faint or undetected but are easily visible in the near-infrared. Detailed classification of the sources with near-infrared spectra is difficult since many sources have featureless continua, consistent with the line photons being obscured / destroyed by large column density gas and the associated dust. The \\c\\ sources are harboured in massive evolved host galaxies, as is expected from the bulge / black-hole inter-related growth which has been inferred in recent years. Thus, through observations of the dominant flux contributors to the X-ray background, we are beginning to find some of the truly exotic accretion sources in the distant Universe."
    },
    "0201/astro-ph0201069_arXiv.txt": {
        "abstract": "A field of 1.013 ster in the ROSAT all-sky survey (RASS), centered on the south galactic pole (SGP), has been searched in a systematic, objective manner for clusters of galaxies. The procedure relied on a correlation of the X-ray positions and properties of ROSAT sources in the field with the distribution of galaxies in the COSMOS digitised data base, which was obtained from scanning the plates of the UK Schmidt IIIa-J optical survey of the southern sky. The study used the second ROSAT survey data base (RASS-2) and included several optical observing campaigns to measure cluster redshifts. The search, which is a precursor to the larger REFLEX survey encompassing the whole southern sky, reached the detection limits of both the RASS and the COSMOS data, and yielded a catalog of 186 clusters in which the lowest flux is $1.5\\times 10^{-12}\\,{\\rm erg}\\,{\\rm cm}^{-2}\\,{\\rm s}^{-1}$ in the $0.1-2.4\\,{\\rm keV}$ band. Of these 157 have measured redshifts. Using a flux limit of $3.0\\times 10^{-12}\\,{\\rm erg}\\,{\\rm cm}^{-2}\\,{\\rm s}^{-1}$ a complete subset of 112 clusters was obtained, of which 110 have measured redshifts. The spatial distribution of the X-ray clusters out to a redshift of $0.15$ shows an extension of the Local Supercluster to the Pisces-Cetus supercluster $({\\rm z}<\\sim 0.07)$, and an orthogonal structure at higher redshift $(0.07<{\\rm z}<0.15)$. This result is consistent with large-scale structure suggested by optical surveys. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0201/astro-ph0201319_arXiv.txt": {
        "abstract": "We observed the classical T Tauri star TW Hya with \\textit{HST}/STIS using the E140M grating, from 1150--1700 \\AA, with the E230M grating, from 2200--2900 \\AA, and with \\FUSE\\ from 900--1180 \\AA.  Emission in 143 Lyman-band H$_2$ lines representing 19 progressions dominates the spectral region from 1250--1650 \\AA. The total H$_2$ emission line flux is $1.94 \\times 10^{-12}$ erg cm$^{-2}$ s$^{-1}$, which corresponds to $1.90\\times10^{-4}$ $L_\\odot$ at TW Hya's distance of 56 pc.  A broad stellar \\Lya\\ line photoexcites the H$_2$ from excited rovibrational levels of the ground electronic state to excited electronic states. The \\ion{C}{2} 1335 \\AA\\ doublet, \\ion{C}{3} 1175 \\AA\\ multiplet, and \\ion{C}{4} 1550 \\AA\\ doublet also electronically excite H$_2$. The velocity shift of the H$_2$ lines is consistent with the photospheric radial velocity of TW Hya, and the emission is not spatially extended beyond the $0\\farcs05$ resolution of \\textit{HST}. The H$_2$ lines have an intrinsic FWHM of $11.91\\pm0.16$ \\kms. One H$_2$ line is significantly weaker than predicted by this model because of \\ion{C}{2} wind absorption.  We also do not observe any H$_2$ absorption against the stellar \\Lya\\ profile.  From these results, we conclude that the H$_2$ emission is more consistent with an origin in a disk rather than in an outflow or circumstellar shell. We also analyze the hot accretion-region lines (e.g., \\ion{C}{4}, \\ion{Si}{4}, \\ion{O}{6}) of TW Hya, which are formed at the accretion shock, and discuss some reasons why Si lines appear significantly weaker than other TR region lines. ",
        "introduction": "T Tauri stars (TTSs) are young ($< 10$ Myr), roughly solar mass stars that have only recently emerged from their natal molecular clouds to become optically visible \\citep[see reviews by][]{Ber89,Fei99}.  The presence of disks around classical TTSs (CTTSs) is now well established observationally.  Millimeter wave measurements of the dust continuum from the outer disk can be used to estimate the disk mass, assuming a dust to gas ratio (see review by Zuckerman 2001).  Such estimates typically yield masses of $\\sim 0.02$ M$_\\odot$, consistent with the minimum mass for the solar nebula \\citep[e.g.,][]{Bec90}.  Spectral line observations can be used to study the inner regions of the disk through their kinematic signatures.  This has been done successfully for very young embedded sources using the infrared (IR) bands of CO \\citep*{Car01,Thi01,Naj96,Cha95}, but the results are confusing for the CTTSs that show CO band emission \\citep{Thi01,Cha95}.  Molecular hydrogen is expected to be $\\sim 10^4$ times more abundant than other gas tracers such as CO in the disks surrounding young stars.  Depending on the density, H$_2$ can survive at temperatures of up to 4000 K \\citep[e.g.,][]{Wil00a}, and can self-shield well against UV radiation fields, making it an excellent diagnostic of disks around young stars.  Good probes of gas in circumstellar disks are needed to determine the lifetime of gas in these environments where planets likely form.  Most studies of disks to date have relied on the IR to millimeter wavelength spectral energy distribution (SED) of these young stars.  These studies are mainly sensitive to micron-sized dust and have shown that the typical survival time for this dust is only a few million years \\citep*{Str93,Wol96,Hil99,Alv00}.  This does not mean, however, that the disks around these stars disappear on this short timescale.  Theories of giant planet formation suggest that after the dust collects into larger particles (and no longer shows up in studies of the SED), the gas is still present in the disk for some time before it accretes onto the planets \\citep*[e.g.,][]{Wuc00}.  Indeed, if for most CTTSs the gas disappears from the disk on the same timescale as the dust does, it may be quite difficult to form giant gas planets.  Molecular hydrogen around a CTTS was first detected by 2.12183 $\\mu$m 1-0 S(1) line emission in the spectrum of T Tau \\citep{Bec78}. \\citet{Bro81} discovered ultraviolet emission lines of H$_2$ in an \\IUE\\ spectrum of T Tau and identified them as fluorescent lines pumped by \\Lya, which were previously seen in sunspot spectra.  In both cases, it was apparent that the molecular hydrogen emission from T Tau is spatially extended.  Subsequently, H$_2$ has been studied often around CTTSs in both the IR and UV. \\citet{Car90} surveyed young stellar objects in the IR, finding H$_2$ emission in 4 CTTSs.  The location of this H$_2$ emission is very important.  Recent ground-based images of T Tau show that the IR emission is quite extended, reaching to 20\\arcsec\\ from the star \\citetext{\\citealp{Van94}; see also \\citealp*{Her96,Her97}}.  The extended IR H$_2$ emission seen around T Tau and other young embedded sources like L1448 \\citep*{Bal93} is interpreted as shock heated emission from the interaction of the stellar/disk wind/jet with the surrounding cloud material \\citep[see review by][]{Rei01}. \\citet{Thi01} used ISO to detect H$_2$ emission in the IR from 22 CTTSs, Herbig Ae/Be stars, and debris-disk stars, but with a large aperture extending over 20\\arcsec\\ this emission may include both disk emission and spatially extended H$_2$ emission from outflows.  They assumed based on the strength of on-source versus off-source CO emission that the bulk of the emission occurred in the circumstellar disks.  However, \\citet{Ric01} attemped to detect on-source H$_2$ emission from five of these sources using a small slit, but failed to detect any significant H$_2$ emission.  Magnetic quadrupole transitions of H$_2$ in the IR are very weak, and electric dipole transitions within the ground electronic state are forbidden because H$_2$ has no net dipole moment.  In the UV, the large oscillator strengths of the electronic Lyman-band (B-X) transitions lead to strong emission and absorption lines via transitions to and from the excited electronic state. \\citet*{Val00} observed H$_2$ emission in the UV around 13 of 32 TTSs studied with \\IUE.  \\citet{Val00}  go on to argue that the limited sensitivity of \\IUE\\ may be the only reason that H$_2$ was not detected in the remaining 19 TTS.  \\citet*{Joh00} attributed UV pumping to the excitation of H$_2$, because the observed flux in H$_2$ lines scales with the total UV flux. The UV H$_2$ lines of T Tau are spatially extended beyond the disks thought to be in this system and may have substantial contributions from outflows.  However, as described below, we do not think that the H$_2$ lines from TW Hya arise in a molecular outflow.  More recently, \\citet{Ard01} detected fluorescent H$_2$ around a sample of CTTSs with \\HST/GHRS and attributed the pumping to the red wing of \\Lya.  \\citet{Rob01} and \\citet{Her01} presented observations of circumstellar H$_2$ absorption against the \\ion{O}{6} profiles of the Herbig Ae/Be stars AB Aur and DX Cha, as observed with the Far Ultraviolet Spectrograph Explorer (\\FUSE).  \\citet{Her01} identified H$_2$ emission in \\FUSE\\ observations of TW Hya and V4046 Sgr in two Lyman-band H$_2$ lines, but did not detect significant H$_2$ absorption against \\ion{O}{6} for these two stars.  TW Hya is the most prominent member of the recently named TW Hydrae association \\citep[TWA --][]{Web99} located at a distance of $56\\pm8$ pc \\citep{Wic98}, making it the closest known CTTS to the solar system. \\citet{Web99} estimated that TW Hya is a $0.7$ $M_\\odot$ star with an age of 10 Myr based on evolutionary tracks of \\citet{Dan97}.  Most other members of the TWA are naked T Tauri stars (NTTSs - TTSs which lack close circumstellar disks), which confirms that TW Hya is quite old for a CTTS. TW Hya shows strong H$\\alpha$ emission \\citep{Ruc83,Web99} produced in the accretion column, and an IR excess \\citep{Jay99} produced by dust emission from a disk. The mass accretion rate has been measured to be from $5-100\\times10^{-10}M_\\odot$ yr$^{-1}$ \\citep{Muz00,Ale01}. By measuring the width of magnetically sensitive lines, \\citet{Joh01a} estimated that the mean magnetic field of TW Hya is of order 3 kG. Recent imaging of TW Hya shows that its disk is viewed very nearly face-on \\citep{Wei99,Kri00}.  D. Potter \\citetext{2001, private communication} determined the inclination angle $i=10\\pm5^\\circ$ using IR polarimetry, while \\citet{Ale01} estimated an inclination of $i=17.5\\pm4.5^\\circ$ by analyzing time-series profiles of H$\\alpha$. The outer radius of the disk extends more than 225 AU \\citep{Kri00,Tri01,Wei01} from the star, and its inner truncation radius is about 0.3 AU \\citep{Tri01}.  \\citet{Dal01} found evidence from comparing the observed spectral energy distribution (SED) of TW Hya to models of SEDs that grains in the inner $0.5$ AU of the disk may have grown in size and settled towards the midplane, while dust and gas in the outer disk is well-mixed.  The disk mass is roughly $1.5-3 \\times 10^{-2} M_\\odot$ \\citep{Wil00,Thi01,Tri01}. Based on submillimeter CO emission and estimates for the disk mass, \\citet{Thi01} and \\citet{Zad01} determined that CO is underabundant by a factor of $\\sim500$ around TW Hya, with respect to a typical ISM H$_2$/CO ratio of $10^4$. There is no molecular cloud associated with TW Hya \\citep{Web99}, and therefore no opportunity to observe whether or not there is an outflow from H$_2$ or CO observations.  \\citet*{Wei00} reported the detection of H$_2$ emission in the 1-0 S(1) line at 2.12183 $\\mu$m from TW Hya, attributing this line to formation in the circumstellar disk, in which the H$_2$ is excited by non-thermal electrons produced by X-ray ionization. TW Hya is a strong X-ray source, with $L_x\\sim 4\\times 10^{-4}$ $L_\\odot$ as measured with Chandra by \\citet{Kas01}. These Chandra observations indicated an overabundance of Ne, an underabundance of O and Fe, and solar-like abundances of Si, Mg and N in the X-ray emission region.  From the sharply peaked emission measure distribution and high electron density, \\citet{Kas01} concluded that the X-ray emission is produced in an accretion shock.  Alternatively, \\citet{Gag02} argued that the Chandra X-ray emission may be consistent with an origin either in the accretion shock or in a corona.  Here, we use \\HST/STIS\\ and \\FUSE\\ spectra of TW Hya to study its circumstellar environment, focusing on the substantial number of H$_2$ emission lines seen in the UV.  In \\S\\ref{section:observations} we explain the observations and data reduction.  In \\S\\ref{section:analysis} we identify the emission lines and examine their profiles.  In \\S4 we show that the Si lines are anonomously weak, and propose a depletion of Si in the accretion column as one possible explanation.  In \\S5, we discuss the H$_2$ fluorescence, and in \\S6 we argue that this emission most likely originates in the circumstellar disk of TW Hya. In paper II, we model the H$_2$ line fluxes to assess the physical conditions that give rise to the H$_2$ fluorescence. ",
        "conclusions": "  1.  The 1250--1650 \\AA\\ spectrum is dominated by molecular hydrogen fluorescence in 143 Lyman band lines.  The observed flux in these lines is $1.94 \\times 10^{-12}$ erg cm$^{-2}$ s$^{-1}$, corresponding to $1.90 \\times 10^{-4} L_\\odot$ at the stellar distance of 56 pc.  The H$_2$ spectrum is pumped to the excited electronic state mainly by transitions coincident with the broad \\Lya\\ emission line.  Other lines, such as the \\ion{C}{2} 1335 \\AA\\ doublet, \\ion{C}{3} 1175 \\AA\\ multiplet, and the \\ion{C}{4} 1550 \\AA\\ doublet may also pump some H$_2$ to the excited electronic state.  2.  The H$_2$ emission is not spatially extended from the star, given a resolution of $0\\farcs05$, corresponding to 2.8 AU at 56 pc.  With a face-on disk, the H$_2$ emission is produced within 1.5 AU of the central star.   3.  The wavelengths of the H$_2$ emission is not shifted by more than 3 \\kms\\ with respect to the photospheric velocity of the star.  4.  The stellar wind occurs in our line of sight to the H$_2$ emission region, based upon the anomolously weak flux in an H$_2$ line that has a wavelength coincident with a wind feature of \\ion{C}{2} 1334.5 \\AA.  5.  No significant H$_2$ absorption is detected against the \\Lya\\ profile or the \\ion{O}{6} emission lines.  6.  The H$_2$ emission likely originates in the surface layer of a disk, rather than a circumstellar shell or an outflow.  We cannot rule out the possibility that the H$_2$ emission is produced in the stellar photosphere.  7.  Silicon II, III, and IV lines are anomolously weak in the UV, possibly because Si depletes onto grains that accumulate in the disk midplane and decouples from the surface material of the disk that preferentially accretes.  Thus the accretion column responsible for the high temperature emission lines would then be very silicon-poor.  Models of the accretion shock are needed to further probe this possibility."
    },
    "0201/astro-ph0201497_arXiv.txt": {
        "abstract": "In the context of star formation through fragmentation of an extremely metal-deficient protogalactic cloud, the gravitational collapse of filamentary gas clouds is explored with one-dimensional numerical hydrodynamics coupled with non-equilibrium chemistry of H$_2$ and HD. It is found that the cloud evolution is governed mainly by the initial central density ($n_{\\rm c,0}$) and H$_2$ abundance ($x_{\\rm H_2,0}$). In particular, the evolution of low-density filaments ($n_{\\rm c,0} \\lesssim 10^5$ cm$^{-3}$) bifurcates at a threshold H$_2$ abundance of $x_{\\rm H_2,cr}\\simeq 3\\times 10^{-3}$, beyond which HD cooling overwhelms H$_2$ cooling. The contraction of a filament with $n_{\\rm c,0} \\lesssim 10^5$ cm$^{-3}$ and $x_{\\rm H_2, 0}\\gtrsim x_{\\rm H_2,cr}$ is strongly decelerated when the central density ($n_{\\rm c}$) reaches a critical density of HD at which LTE level populations are achieved, and therefore the filament is expected to fragment at $\\sim 10^{7}$ cm$^{-3}$. The fragment mass is assessed to be $\\approx 10M_\\odot$. In contrast, the contraction of a filament with $n_{\\rm c,0} \\lesssim 10^5$ cm$^{-3}$ and $x_{\\rm H_2, 0}\\lesssim x_{\\rm H_2,cr}$ is regulated by H$_2$ cooling. In this case, the filament tends to fragment at lower density as $\\sim 10^{4}$ cm$^{-3}$ owing to the low critical density of H$_2$, and the fragment mass is as high as $\\approx 10^2M_\\odot$. For a high-density filament with $n_{\\rm c,0} \\gtrsim 10^5$ cm$^{-3}$, the temperature stays at a relatively high value because both H$_2$ and HD cooling saturate, and the cloud evolution is governed by H$_2$ cooling. The contraction of a high-density filament is accelerated by effective three-body H$_2$ formation when the density reaches $10^{8-9}$ cm$^{-3}$. The fragmentation is not expected to take place until the cloud becomes opaque in H$_2$ lines at $n_{\\rm c,0}\\sim 10^{12-13}$ cm$^{-3}$, so that the fragment mass is reduced to $1-2$ M$_\\odot$. As a result, the stellar initial mass function (IMF) could be bimodal and deficient in sub-solar mass stars, where the high mass peak is around $10M_\\odot$ or $10^2M_\\odot$, dependently on $n_{\\rm c,0}$ and $x_{\\rm H_2,0}$. If the protogalactic clouds are ionized by UV radiation or strong shocks, the H$_2$ abundance could exceed $x_{\\rm H_2,cr}\\simeq 3\\times 10^{-3}$ by reactions of $\\rm H+e \\rightarrow H^- + {\\it h\\nu}$ and $\\rm H+H^-\\rightarrow H_2+e$. The high mass peak would then be $O(10) M_\\odot$. ",
        "introduction": "The first generation of stars, Population III stars, are thought to form from almost metal-free gas with metallicity of $\\sim 10^{-10} Z_\\odot$ (see e.g., Carr, Bond, \\& Arnett 1984; Carr 1994). Such Population III stars may have produced heavy elements at redshifts $z \\gtrsim 10$. Recent observations of quasar absorption spectra show that the metal abundance in intergalactic space is $\\approx 10^{-3} Z_\\odot$ (e.g., Cowie \\& Songaila 1998). Also, it is widely accepted that a significant portion of old halo stars in our Galaxy have metallicity lower than $10^{-3}Z_\\odot$ (Beers, Preston, \\& Shectman 1992). In addition, some blue compact dwarf galaxies are known to be extremely metal-poor ($\\approx 10^{-2} Z_\\odot$, Kunth \\& Sargent 1986; Pustilnik et al. 2001). Thus, star formation in protogalaxies must have proceeded in very metal-deficient environments.   When the metallicity is lower than $Z \\lesssim 10^{-2}Z_\\odot$, the cooling by heavy elements is less effective than cooling by hydrogen and helium (e.g., Yoshii \\& Sabano 1980; B\\\"ohringer \\& Hensler 1989; Omukai 2000; Nishi \\& Tashiro 2000). For the formation of Population III stars from such metal-deficient gas, the cooling by primordial hydrogen molecules (H$_2$) plays an essential role (Matsuda, Sato, \\& Takeda 1969; Yoneyama 1972; Hutchins 1976; Silk 1977; Yoshii \\& Sabano 1980; Carlberg 1981; Lepp \\& Shull 1984; Palla, Salpeter, \\& Stahler 1983; Yoshii \\& Saio 1986; Shapiro \\& Kang 1987; Uehara et al. 1996; Haiman, Thoul \\& Loeb 1996; Nishi et al. 1998; Abel et al. 1998; Bromm, Coppi, \\& Larson 1999; Abel, Bryan, \\& Norman 2000; Nakamura \\& Umemura 1999a, 2001, hereafter Papers I and II). Recent studies have revealed that star formation in primordial gas is considerably different from present-day star formation, implying that the stellar initial mass function (IMF) might be different from the Salpeter-like IMF and the IMF might be time-varying in the course of galaxy evolution from the early collapsing stages to the present day (see also Larson 1998; Zepf \\& Silk 1996; Chabrier 1999).  Besides hydrogen molecules, deuterated hydrogen molecules (HD) can be a significant coolant (e.g., Galli \\& Palla 1998). Although HD is less abundant than H$_2$ ([HD/H$_2$] $\\sim 10^{-3}- 10^{-4}$), HD has a finite dipole moment and thus higher radiative transition probabilities than H$_2$. For example, the lowest rotational transitions of H$_2$ and HD have radiative transition probabilities of $A_{20}=3\\times 10^{-11}$ s$^{-1}$ and $A_{10}=5\\times 10^{-8}$ s$^{-1}$, respectively. Also, the corresponding excitation energies of H$_2$ and HD are $\\Delta E_{20}/k = 510$ K and $\\Delta E_{10}/k = 128$ K, respectively, where $k$ is the Boltzmann coefficient. Thus, HD can lower the gas temperature down to $T\\lesssim$ 100K (e.g., Puy \\& Signore 1996; Bougleux \\& Galli 1997; Galli \\& Palla 1998; Flower et al. 2000).  The first pregalactic objects should have collapsed at redshifts of $z\\sim 10 - 10^2$ and have masses of $10^5 - 10^8$ M$_\\odot$ in a cold dark matter cosmology (Tegmark et al. 1998; Fuller \\& Couchman 2000). When such first objects are virialized and the gas temperature ascends to $10^3 - 10^4$ K, H$_2$ formation is promoted and the fractional abundance is raised to $x_{\\rm H_2}=10^{-4} - 10^{-3}$. However, for this H$_2$ abundance, the H$_2$ cooling cannot lower the temperature down to 100 K. Thus, the cloud evolution is basically controlled by H$_2$ cooling and HD cooling is not important (Nakamura \\& Umemura 1999b, 2000; see also Lepp \\& Shull 1984).  But, in some situations, the H$_2$ abundance can be raised further. Ferrara (1998) argued that in dense shells formed behind supernova shocks, H$_2$ molecules form efficiently owing to reduced recombination of free electrons (see also Shapiro \\& Kang 1987). In the dense shells, the H$_2$ concentration reaches $6 \\times 10^{-3}$, and consequently the gas temperature descends to $\\sim 100$ K. At such low temperature, HD is a more efficient coolant than H$_2$. The primordial star formation regulated by HD cooling in shocked shells is studied by several authors (Uehara \\& Inutsuka 2000; Machida, Fujimoto, \\& Nakamura 2001). Another possibility is star formation from photoionized gas. Corbelli, Galli, \\& Palla (1998) studied the effects of UV background radiation on the thermal evolution of the protogalaxies (see also Murray \\& Lin 1989; Haiman, Rees, \\& Loeb 1996). They found that at redshifts lower than $z \\sim 1-2$, protogalactic clouds are self-shielded against the UV background radiation, and H$_2$ formation is promoted up to $x_{\\rm H_2} \\approx 10^{-2}$ with the help of abundant free electrons due to the UV radiation. Even at higher redshifts ($z\\gtrsim 2$), if pregalactic clouds more massive than $10^{11} M_\\odot$ collapse under a UV background, they can be eventually self-shielded and abundant H$_2$ ($\\gtrsim 10^{-3}$) forms, so that the temperature descends down to $\\sim 100$K (Susa \\& Umemura 2000). Hence, HD cooling is likely to be important for the star formation in the protogalactic systems. In this paper, we extensively examine the effects of the HD cooling on the formation of stars in protogalaxies, and elucidate the role of HD molecules for the stellar initial mass function (IMF) there.  This paper is organized as follows.  In \\S 2, we describe model and numerical methods which are basically the same as those of Paper II.  Numerical results are presented in \\S 3. In \\S 4, implications for the IMF of Population III and metal-deficient stars are discussed. ",
        "conclusions": "First, we briefly discuss the role of HD for the formation of the very first stars (Pop III stars). In the bottom-up scenarios like cold dark matter models, when the first pregalactic objects with masses of $10^5 - 10^8$ M$_\\odot$ are virialized at redshifts of $z\\sim 10 - 10^2$, the H$_2$ abundance reaches at most $10^{-4} - 10^{-3}$ which is lower than the threshold H$_2$ abundance. Therefore, the cloud evolution is basically determined by H$_2$ cooling rather than HD cooling (Nakamura \\& Umemura 1999b, 2000). If the primordial D abundance is a few times as high as the value we assumed ($x_{\\rm D}=4\\times 10^{-5}$), HD cooling may play a role in the thermal evolution of the pregalactic clouds. However, recent observations of quasar absorption spectra find an observed primordial D abundance as low as $3-4\\times 10^{-5}$ (Tytler et al. 1996), implying that HD cooling is not important for the first star formation.   Next, we give some implications for star formation in metal-deficient primordial galaxies. Recently, Susa \\& Umemura (2000) investigated the pancake collapse of pregalactic clouds under UV background radiation. They found that once the pancake disk is shielded against external UV radiation in the course of contraction, H$_2$ molecules form efficiently via the H$^-$ reaction with abundant free electrons produced by UV background, and the resultant abundance reaches $x_{\\rm H_2} \\approx 3\\times 10^{-3}$ (see also Shapiro \\& Kang 1986). The pancake disks probably fragment into filaments in which stars can subsequently form. In this case, HD cooling is expected to become efficient in low-density filaments, and then the high mass peak of the IMF would go down to $\\sim$ 10 $M_\\odot$.   It is found that at redshifts of $z\\sim 2$, the UV background radiation decreases with time (Irwin, McMahon, \\& Hazard 1991; Maloney 1993). The time-decreasing UV background radiation is likely to influence star formation in galaxies, especially low surface brightness galaxies (e.g., Ellis 1997). Corbelli et al. (1998) studied the effects of the declined UV background radiation on the thermal evolution of the protogalaxies with low surface densities. They found that there is a critical redshift of $z \\sim 1-2$, above which the gas disks with surface densities $10^{20}$ cm$^{-2} \\lesssim N_{\\rm HI} \\lesssim 10^{21}$ cm$^{-2}$ are gravitationally stable at $T\\sim 10^4$ K. Below this redshift, the declined UV radiation is shielded by the gas disks where the H$_2$ abundance reaches $10^{-2}$ owing to high ionization degree by the UV radiation. Also, in such galaxies, the high mass peak of the IMF would decrease to $\\sim 10M_\\odot$ owing to the HD cooling.  The high mass end of the IMF can influence abundance patterns because metal production by extremely metal-deficient stars is very different between $10^2 M_\\odot$ and $10 M_\\odot$ (Abia et al. 2001; Umeda \\& Nomoto 2001; Heger \\& Woosley 2001; Schaerer 2001). For instance, the abundance pattern in the metal-poor ($\\sim 0.05 Z_\\odot$) starburst galaxy M82 cannot be accounted for unless stars with $\\gtrsim 25 M_\\odot$ contribute  significantly to the metal enrichment of the galaxy (Tsuru et al. 1997; Nakamura et al. 2001; Umeda \\& Nomoto 2001). This mass scale seems to be consistent with the high mass peak of the present IMF regulated by the HD cooling. It is also consistent with the estimate by Hernandez \\& Ferrara (2001).  As the star formation progresses, the interstellar metallicity will monotonously increase with time.  When the metallicity reaches $10^{-3} - 10^{-2}Z_\\odot$, the metal cooling becomes important and the thermal properties of the gas are changed. Thereafter, the process of star formation would become similar to the present-day case. In other words, the IMF would settle into Salpeter-like IMF."
    },
    "0201/astro-ph0201174_arXiv.txt": {
        "abstract": "{ It is shown that linear polarization data can be used to constrain the composition (normal or pair plasma) of pc-scale extragalactic jets. A simple criterion, based on synchrotron and Faraday depolarization properties, is established. It does not depend on the particle density and the length of the emitting region along the line of sight, thus eliminating two physical unknowns. ",
        "introduction": "Due to their synchrotron emission, extragalactic jets are known to contain electrons. But what is the positive charge population : protons or positrons? This makes a huge difference in the kinetic energy to be put in the ejection process and implies different properties for acceleration and propagation mechanisms.  Extracted matter from the accretion disk is comprised of normal (electron-proton) plasma. But accelerating protons to relativistic speeds makes the energetic budget somewhat problematic. The scheme of a two--fluid plasma (Sol et al. \\cite{Sol}, Pelletier \\& Marcowith \\cite{pelletier}) alleviates this difficulty by reserving the relativistic speed to a pair (electron--positron) plasma, while the bulk of the jet, non relativistic and non radiating, is made of normal plasma.  Observationally, there have been two approaches to determine the nature of the radiating plasma: spectroscopic behavior (Reynolds et al. \\cite{Reynolds}), and circular polarization (Wardle et al. \\cite{Wardle}). The results tend very strongly toward electron--positron pairs for the radiating component.  In this paper, it is shown that linear polarization could also be used to derive the nature of the jet material from very high resolution observations. Based on basic formulae presented in Sect.~2, a very simple analytical criterion is established in Sect.~3. In Sect.~4, some particular observational data are discussed before a brief conclusion is given in Sect.~5. ",
        "conclusions": "The present study has been initiated from the synchrotron emission simulations (Despringre \\& Fraix-Burnet \\cite{paper1}, Fraix-Burnet in prep.) in which it appears nearly impossible to reproduce both surface brightness and polarization of typical VLBI jets. We have then devised a simple tool, presented in this paper, to constrain the composition of a jet with linear polarization data.  It is based on the Faraday depolarization which is the minimum depolarization that occurs from a normal plasma. This yields a criterion, easily computed from the observations, that does not depend on the length of the emitting plasma along the line of sight nor on density of the emitting particles. These two parameters are always very difficult to estimate.  Still, there are a few uncertain parameters in the criterion, but from a few illustrative examples presented in this paper, it already appears that they should be pushed to high values if one assumes a normal plasma. A complete statistical study of several jets or a very detailed study of a well-observed jet is necessary to reach more significant conclusions. This is beyond the scope of the present paper.   Our criterion is not valid in the self--absorbed part of the spectrum. It is certainly not a problem for frequencies above 10~GHz in the jets, but it should be applied to cores only with caution."
    },
    "0201/astro-ph0201342_arXiv.txt": {
        "abstract": "If the first (PopIII) stars were very massive, their final fate is to collapse into very massive black holes. Once a proto-black hole has formed into the stellar core, accretion continues through a disk. It is widely accepted, although not confirmed, that magnetic fields drive an energetic jet which produces a burst of TeV neutrinos by photon-meson interaction, and eventually breaks out of the stellar envelope appearing as a Gamma Ray Burst (GRB). Based on recent numerical simulations and neutrino emission models, we predict the expected neutrino diffuse flux from these PopIII GRBs and compare it with the capabilities of present and planned detectors as AMANDA and IceCube. If beamed into 1\\% of the sky, we find that the rate of PopIII GRBs is $\\le 4 \\times 10^6$~yr$^{-1}$. High energy neutrinos from PopIII GRBs could dominate the overall flux in two energy bands [$10^4 - 10^5$] GeV and [$10^5 - 10^6$] GeV of neutrino telescopes. The enhanced sensitivities of forthcoming detectors in the high-energy band (AMANDA-II, IceCube) will provide a fundamental insight on the characteristic explosion energies of PopIII GRBs and will constitute a unique probe of the the Initial Mass Function (IMF)  of the first stars and of the redshift $z_f$ marking the metallicity-driven transition from a top-heavy to a normal IMF. The current upper limit set by AMANDA-B10 implies that such transition must have occurred not later than $z_f =9.2$ for the most plausible jet energies. Based on such results, we speculate that PopIII GRBs, if not chocked, could be associated with a new class of events detected by BeppoSax, the Fast X-ray Transient (FXTs), which are bright X-ray sources, with peak energies in the $2-10$~keV band and durations between $10-200$~s. ",
        "introduction": "One of the most challenging problems in modern cosmology is the understanding of the first episodes of star formation in the Universe. A growing body of theoretical work has been devoted in the past few years to this subject (Rees 1976; Rees \\& Ostriker 1977; Silk 1977, 1983; Haiman, Thoul \\& Loeb 1996; Uehara \\etal 1996). These investigations are aimed at the identification of the characteristic mass scale of the first stars (usually referred to as PopIII stars), how this relates to the physical conditions of the gas in the first collapsed objects and how it affects subsequent structure formation in the Universe.  Numerical simulations (Abel \\etal 1998; Nakamura \\& Umemura 1999; Bromm, Coppi \\& Larson 1999, 2001; Abel, Bryan \\& Norman 2000; Bromm \\etal 2001; Ripamonti \\etal 2001) based on hierarchical scenarios of structure formation and/or detailed stellar collapse models have shown that the typical mass scale for the first collapsed clumps of primordial gas is $\\approx 10^{3} \\msun$, which corresponds to the Jeans mass set by molecular hydrogen cooling.  The ultimate nature of the stars that form out of these clumps critically depends on the physical conditions of the gas as the evolution pushes density to higher values. Preliminary studies (Omukai \\& Nishi 1998; Nakamura \\& Umemura 1999; Bromm \\etal 2001; Ripamonti \\etal 2001) show that as long as the metallicity is below some critical value (typically $Z_{\\rm cr} \\sim 10^{-4} Z_{\\odot}$) the first clumps have little tendency to further fragment, as also expected from a number of physical arguments (see \\eg Schneider \\etal 2001).  Tentative evidences for an early top-heavy initial mass function (IMF) are provided by a number of observations (see Hernandez \\& Ferrara 2000 and references therein). For instance, a comparison between the observed number of metal poor stars with the one predicted by cosmological models, implies that the stellar characteristic mass sharply increases with redshift. Furthermore, the intracluster medium (ICM) metal abundances measured from {\\it Chandra} and {\\it XMM} spectral data are higher than expected from the enrichment by standard IMF supernova (SN) yields in cluster galaxy members, indicative of a top-heavy early IMF.  Finally, the observed abundance anomalies (\\eg oxygen) in the ICM can be explained by an early generation of PopIII SNe (Loewenstein 2001).  These issues, highly suggestive of a top-heavy early star formation, have recently motivated a series of numerical investigations of the nucleosynthesis and final fate of metal-free massive stars (Heger \\& Woosley 2001; Fryer, Woosley \\& Heger 2001). Stars with masses in the range $140 \\msun \\leq M \\leq 260 \\msun$ undergo electron-positron pair instability  (\\SNgg) and end up in a giant, nuclear-powered explosion, leaving no remnant and enriching the ambient medium with their nucleosynthetic products. The kinetic energy released during the thermonuclear explosions powered by pair instability are $\\approx 10^2$ times larger than those of ordinary Type II SNe. This might cause the interaction with the circumstellar medium to be as strong as predicted for hypernovae (Woosley \\& Weaver 1982). These explosions do not lead to the ejection of strongly relativistic matter and therefore cannot power a Gamma-Ray Burst (GRB, Fryer, Woosley \\& Heger 2001).  However, if stars have masses $M>260 \\msun$, photodisintegration instability is encountered before explosive nuclear burning can reverse the implosion and the stars collapse to Very Massive Black Holes (VMBHs), swallowing virtually all previously produced heavy elements. These stars are likely to be rapidly rotating and the estimated angular momentum is sufficient to delay black hole formation (Fryer, Woosley \\& Heger 2001). Once a proto-black hole has formed into the core, accretion continues through a disk at a rate which can be as large as $1-10\\, \\msun \\, {\\rm s}^{-1}$. It is widely accepted, although not confirmed, that magnetic fields might drive an energetic jet which can produce a strong GRB through the interaction with surrounding gas.  In this scenario, the energetic jets generated by GRBs engines produce, by photon-meson interaction, a burst of TeV neutrinos while propagating in the stellar envelope (M{\\'e}sz{\\'a}ros \\& Waxman 2001). It is widely assumed that if the progenitors of GRBs are massive, collapsing stars (the so called ``collapsar'', Woosley 1993, Paczy\\'nski 1998) the shocks producing the $\\gamma$-rays occur after the relativistic jet has emerged from the stellar envelope. For a significant fraction of collapsars, the jet may be unable to punch through the stellar envelope (MacFadyen, Woosley \\& Heger 2001). However, the TeV neutrino signal from such ``chocked\" jets should be similar to that from jets which do break through the stellar envelope, leading to observable GRBs. Detecting the neutrino flux from individual collapsars may provide a direct evidence of such $\\gamma$-ray-dark collapses (M{\\'e}sz{\\'a}ros \\& Waxman 2001).  In this paper, we compute the diffuse flux of high energy neutrinos emitted by PopIII GRBs at high redshifts ($5.4 < z < 30$). Hereafter, by PopIII GRBs we indicate PopIII stars with masses $>260 \\msun$ which collapse to a VMBHs after a transient phase of mass accretion.  Following Schneider \\etal 2001, we assume that a fraction $(1-\\fgg)$ of the first stars collapse to VMBHs powering a relativistic jet that generates a flux of high energy neutrinos. Using the neutrino emission model proposed by M{\\'e}sz{\\'a}ros \\& Waxman (2001) and integrating over the source rate throughout the universe, we compute the diffuse neutrino flux from PopIII GRBs.  The sensitivities of present (AMANDA\\footnote{\\tt http://amanda.berkeley.edu/amanda/amanda.html}, Andr{\\'e}s \\etal 2001) and forthcoming neutrino telescopes (such as AMANDA-II, Barwick 2001, IceCube\\footnote{\\tt http://www.ssec.wisc.edu/a3ri/icecube/}, Spiering 2001 and other km-scale detectors, Halzen 2001) enable to derive important constraints on the first episodes of star formation, such as the relative number of \\SNgg \\, and VMBHs (\\ie PopIII GRBs), metal enrichment and the nature of dark matter.  The paper is organized as follows: in Section 2 we describe the proposed scenario for the formation of PopIII stars and we compute the expected rate of GRBs. In Section 3 we determine the energy spectrum of TeV neutrinos emitted in individual PopIII GRBs. In Section 4 we compute the expected diffuse flux from PopIII GRBs. In Section 5 we explore the parameter space making detailed predictions for the detectability with present and forthcoming neutrino telescopes. In Section 6 we illustrate the implications of the proposed scenario for PopIII GRBs. Finally, in Section 7 we summarize our main results. ",
        "conclusions": "The main results of this paper can be summarized as follows:  $\\bullet$ The total rate of PopIII GRBs produced by massive, first stars can be as high as $4 \\times 10^6$~yr$^{-1}$ for the most favourable parameter choice and if the jets are beamed into 1\\% of the sky; for some model parameters, the event duty cycle is found to be smaller than unity implying that the diffuse neutrino signal might be characterized by a sequence of individual bursts.  $\\bullet$ High energy neutrinos from PopIII GRBs could dominate (above other plausible neutrino sources) the flux in the L [$10^4-10^5$] GeV and H [$10^5 - 10^6$] GeV energy bands which can be explored by present (AMANDA-I) and forthcoming (AMANDA-II, IceCube) neutrino telescopes (Fig.~\\ref{fig02}).  $\\bullet$ Constraints on $\\fgg$ (\\ie on the mass fraction of PopIII stars which ends up as SN$_{\\gamma\\gamma}$) from these experiments will also provide a tremendous insight into the IMF of the first stars, and in particular, on the redshift $z_f$, which marks the metallicity-driven transition from a top-heavy IMF to a the current, standard (\\ie Salpeter-like) one.  $\\bullet$ For explosion energies in the range $E_{\\rm iso} = 6.4 \\times 10^{54} -10^{56}$~erg, the current upper limit established by AMANDA-I in the L- and H-bands can be used to reject a region of the model parameter space, namely $\\fgg \\leq [8 \\times 10^{-3}-0.27]$ corresponding to transition redshifts $z_f \\leq [5.4-9.2]$. Thus, either PopIII GRBs are characterized by smaller explosion energies or massive primordial star formation occurred only at very high redshifts and metals released by \\SNgg ~~enriched the IGM to the critical level $<Z> = 10^{-4} Z_{\\odot}$ before redshift 9.2.   $\\bullet$ The enhanced sensitivities of forthcoming experiments such as AMANDA-II and, in particular, km-scale detectors such as IceCube, might be able to discriminate between the two above possibilities: observations in the H-band will explore a wide range of explosion energies, $E_{\\rm iso}=7.6 \\times 10^{52}-10^{56}$~erg, values for $\\fgg$ in the range $8 \\times 10^{-3} \\le \\fgg \\le 0.98$ and transition redshifts $5.4 \\le z_f \\le 10$.  $\\bullet$ For a subset of model parameters, neutrino emission from individual bursts might be detected with km-scale detectors in the L-band over roughly a year of observations. To be statistically significant, single burst detection should occurr in coincidence with observations of the GRB itself, if the jet is not chocked and can successfully propagate out of the stellar envelope.  $\\bullet$ We have speculated that PopIII GRBs, if they are not chocked, could be associated with a new class of events, the Fast X-ray Transient (FXTs), which are bright X-ray sources, with peak energies in the $2-10$~keV band and durations between $10-200$~s, which are not triggered and not detected in the $\\gamma$-ray range $40-700$~keV. An additional hint that they could be high redshift sources comes from the fact that all FXTs observed so far fall into the GHOST (GRB Hiding Optical Source Transient) category, an occurrence highly suggestive on the absorbing effects of the predominantly neutral intergalactic medium present at epochs prior to cosmic reionization."
    },
    "0201/astro-ph0201032_arXiv.txt": {
        "abstract": "{We have  detected non-Gaussian signatures in the VIRMOS-DESCART weak lensing survey from a measurement of the three-point shear correlation function, following the method developped by Bernardeau, van Waerbeke and Mellier (2002). We obtain a 2.4$\\sigma$ signal over four independent angular bins, or equivalently, a 4.9-$\\sigma$ confidence level detection with respect to measurements errors on scale of about $2$ to $4$ arc-minutes. Both amplitude and shape are found to agree with theoretical expectations that have been investigated for three cosmological models. This supports the idea that the measure corresponds to a  cosmological signal due to the gravitational instability dynamics. Its properties could be used to put constraints on the cosmological parameters, in particular on the density parameter of the Universe, but the error level as well as the cosmic variance are  still too large to permit secure conclusions. ",
        "introduction": "The large-scale structure of the Universe are expected to form from the gravitational growth of initial density perturbations obeying Gaussian statistics. As the Universe expands and the perturbations grow, non-Gaussian features are expected to emerge in the density field due to gravitational dynamics. These features can be  characterized with perturbation theory calculations, which allow to compute for instance the skewness, third moment of the local density probability distribution function (Peebles 1980, Fry 1984, Bernardeau 1992). The reduced skewness of the density field has been found to be quite insensitive to the variance and the cosmological parameters, $\\Omega_m$ (Juszkiewicz et al. 1992, Bernardeau 1994). In contrast,  weak lensing surveys   are sensitive to $\\Omega_m$ since they trace the integrated mass along the line-of-sight which is roughly proportional to the density parameter of the Universe. Weak lensing by large scale structures has been measured by several teams as a coherent distortion field of distant galaxies over large angular distances (Bacon et al. 2000 and 2002; H{\\\"a}mmerle et al.; 2002; Hoekstra et al. 2002; Kaiser et al. 2000; Maoli et al. 2001; R{\\'e}fr{\\'e}gier et al 2002;  Rhodes et al. 2001; Van Waerbeke et al. 2000, 2001 and 2002; Wittman et al. 2000). The projected mass density reconstructed from the distortion field ($i.e.$ the convergence field) can be used for non-Gaussian signatures searches, as shown in Bernardeau, van Waerbeke \\& Mellier (1997, hereafter BvWM97). This work and further studies (Jain  \\& Seljak 1997; van Waerbeke, Bernardeau \\& Mellier 1999) have shown that the non-Gaussian properties of the convergence field can be used as a probe of the cosmological density parameter, with a weak dependence on the cosmological constant $\\Omega_{\\Lambda}$, provided that the redshift distribution of the sources is known.   However, a straight  application of these theoretical considerations to real data sets turned out to be arduous.  The convergence field has to be recovered from a mass reconstruction process which uses a continuous shear field obtained from a smoothed map of the discrete galaxy ellipticities. Unfortunately, survey topologies are generally complex and are alterated by masked areas due to light scattering, bright stars, comet-like reflections, asteroid/airplane tracks, very bright galaxies, etc.... Since the mask sizes cover a range of scales from a few arc-seconds to two degrees and are strongly anisotropic (for instance bright stars preferentially saturate along CCD columns), results of mass reconstruction in such data sets are not yet reliable. An alternative approach is the aperture mass applied to cosmic shear (Schneider et al. 1998), which allows the measurement of the skewness from the distortion field directly, bypassing the mass reconstruction process. So far, our attempts for measuring the skewness of the aperture mass lead to very noisy and un-significant results. Recently, Bernardeau, van Waerbeke \\& Mellier (2002, hereafter BvWM02)  have proposed a new method using some specific patterns in the shear three-point function. This method has also the advantage to bypass the mass reconstruction process. Despite the complicated shape of the three-point correlation pattern, BvWM02 uncovered a specific geometrical property and demonstrated it can be used for the measurement of non-Gaussian features. Their  detection strategy based on this method has been found to be robust, usable in patchy catalogs, and quite insensitive to the topology of the survey.  In the following we apply this method to the VIRMOS-DESCART weak lensing survey done at the Canada-France-Hawaii Telescope. ",
        "conclusions": "The result shown on Figure \\ref{xidetec} is the first detection of non-Gaussian  features in a cosmic shear survey. The signal is detected  with a 2.4-$\\sigma$ confidence level on 4 independent bins which gives a 4.9-$\\sigma$ global confidence level.  Such a result opens the route to break the degeneracy between  $\\Omega_m$ and $\\sigma_8$ in a way which is independent on assumptions beyond the solely hypothesis that large-scale structures grows from gravitational instability of an initial Gaussian field.  We note that the amplitude of the reduced  3-point correlation function exhibits an angular dependence which is in agreement with theoretical expectations. It supports the interpretation of these results as genuine effects of the gravitational dynamics.  The signal is however still too noisy to provide reliable information on cosmological parameters. Moreover, several obstacles have yet to be overcome. It is first important to understand to which level the measurements are contaminated by systematics. This can be done through consistency checks yet to be invented (a statistic which cancels the signal in a non-trivial way, like the $45$ degrees rotation test for the aperture mass will have to be found for the three-point function). We have already checked with the anisotropy of the stars that our signal is unlikely to be dominated by systematics since it would otherwise exhibits a totally different angular dependence. In the next stages, the scientific interpretation of the three-point function measurements will require a significant improvement over several issues:  -The knowledge of the redshift distribution of the sources is crucial for the third order statistics (see BvWM97);  -The source clustering effect might bias the measurement is a significant way (Bernardeau 1998, Hamana et al. 2002) if the width of the source distribution is too large;  -Intrinsic alignment of galaxies have a completely unknown effect on the non-Gaussian properties of the shear field.  Resolving these issues will require progress from both the theoretical/simulation side and from the observations, which are already on their way.  {"
    },
    "0201/astro-ph0201518_arXiv.txt": {
        "abstract": "Recent observations show that chromospheric and coronal activity in late-M and L dwarfs is much lower than in the earlier M types.  This is particularly surprising, given that the late-M and L dwarfs are comparatively very rapid rotators:  in the early M dwarfs, rapid rotation is associated with high activity levels.  One possibility is that that the drop-off in activity in the late-M's and L's is a result of very high electrical resistivities in their dense, cool and predominantly neutral atmospheres.  We calculate the magnetic field diffusivity in the atmospheres of objects with \\teff in the range 3000 - 1500 K (mid-M to late-L), using the atmospheric structure models of Allard and Hauschildt.  We find that the combination of very low ionization fraction and high density in these atmospheres results in very large resistivities, and thus efficient field diffusion.  While both ambipolar diffusion and Ohmic decay of currents due to ion-electron collisions occur, the primary diffusion effects are due to current decay through collisions of charged particles with neutrals.  Moreover, the latter resistivity is a strong function of both effective temperature and optical depth, increasing rapidly as either \\teff or optical depth decreases.  This has three implications.  One, any magnetic field present is increasingly decoupled from atmospheric fluid motions as one moves from mid-M to L.  In the late-M and L dwarfs, atmospheric motions cannot lead to equilibrium field configurations very different from potential ones.  That is, the magnitude of magnetic stresses generated by atmospheric motions is very small in these objects.  We quantify this effect by a simple Reynolds number calculation.  Two, even if magnetic stresses are easily produced by fluid motions in the hot interior (where the coupling between field and matter is good), their propagation up through the atmosphere will be increasingly hampered by the growing atmospheric resistivity as one moves from mid-M to late L.  Three, these cool dwarfs are expected to be fully convective, with magnetic fields that are generated by a turbulent dynamo.  However, the poor coupling between field and matter in the atmosphere suggests that an efficient dynamo cannot be maintained near the surface, but only at large depths.  Since the turbulent dynamo is thought to produce small-scale fields, burying it at great depths would mean that by the time one reaches the atmosphere, the field strengths would be very weak, and increasingly so with lower \\teff (as the dynamo gets pushed further into the interior).  This would exacerbate the difficulty in producing and transporting magnetic stresses in the atmosphere.  In summary, both the generation and propagation of magnetic stresses are increasingly damped with decreasing \\teff, in these cool dwarfs.  As a result, the magnetic free energy available for the support of a corona, chromosphere and activity becomes smaller and smaller with later type.  This can account for the observed drop in activity from mid-M to L, {\\it assuming} that activity in these dwarfs is magnetically driven.  To check the latter assumption, we estimate the emergent acoustic fluxes in these objects through a Lighthill-Proudman calculation.  While the acoustic fluxes also decrease with decreasing \\teff, they appear inadequate to explain the observed \\hal fluxes in mid-M to L dwarfs.  In the absence of acoustic heating, magnetic heating indeed seems the most viable way of generating activity.  Finally, while our calculations do not address flares in late-M and L dwarfs, we speculate that the latter could be created by buoyant flux tubes that are generated in the interior and rise rapidly through the atmosphere, dissipating their associated currents in the highly resistive upper atmospheric layers. ",
        "introduction": "Recent observations indicate that the ``standard'' rotation-activity connection, observed in stars between roughly F5 to M8.5, breaks down in cooler dwarfs.  While high rotation velocities lead to high chromospheric and coronal activity in the former spectral types, M9 and later dwarfs exhibit very low activity levels in spite of being rapid rotators.  A possible physical basis for this behavior is the topic of this paper.  To begin, we first discuss the usual rotation-activity connection in \\S 1.1.  We then examine the evidence for its breakdown in the later dwarfs, and the suggested reasons for this, in \\S 1.2.  We note that, for the purposes of this paper, only M dwarfs M9 and later will be referred to as late M; all others are called early or mid-M.  \\subsection{``Standard'' Rotation-Activity Connection} A connection between rotation and activity is observed in spectral types F5 to M8.  In these objects, faster rotation corresponds to higher levels of chromospheric and coronal activity (as measured by chromospheric activity indicators such as $\\Lhal/\\Lbol$, and coronal ones such as $\\Lx/\\Lbol$).  This occurs till a cutoff rotational velocity is reached; for objects rotating even faster, the activity saturates at some maximum value.  Simultaneously, rotation also evolves as a function of age and spectral type in stars later than F5: rotational velocity decreases with age in these objects, but the spindown timescale is longer for later types.  For stars in the range F5 to M3, these observations are understood through the following paradigm.  These objects have a radiative core and a convective envelope.  Magnetic fields are expected to be generated through the operation of an $\\alpha\\Omega$ dynamo at the core-envelope interface.  Magnetic stresses are created through the dragging of field lines by photospheric fluid motions; the release of these stresses in the upper atmosphere provides energetic support for a corona and chromosphere and drives activity.  The efficiency of the $\\alpha\\Omega$ dynamo (i.e, the rate of field generation) rises strongly with decreasing Rossby number ($R$), where $R = P/{{\\tau}_c}$, the ratio of the stellar rotation period ($P$) to the convective overturn timescale (${\\tau}_c$).  For a given stellar radius and convective timescale, therefore, larger rotational velocities lead to greater dynamo efficiencies and hence higher levels of activity, as observed.  The phenomenon of saturation is not yet well understood.  One proposal is that saturation occurs when the the rotational velocity is sufficiently high to generate fields that cover the entire stellar surface.  Note that, in this $\\alpha\\Omega$ paradigm, it is not rotational velocity per se, but the Rossby number that is the fundamental parameter.  Indeed, Soderblom et al (1993) find that $\\Lhal/\\Lbol$ in Pleiades G and K dwarfs is better correlated with Rossby number than with \\vsini alone; similar results have been obtained for different samples by other investigators.  Concurrently, for stars between F5 and M3, angular momentum is extracted from the convective envelope and lost through a magnetized wind, leading to a gradual spindown of the star.  As a result, rotational velocity, magnetic field generation and activity all decrease with time.  The convective envelope deepens as one moves to later types.  Therefore the fraction of stellar mass within the convective envelope, and thus the fractional moment of inertia of the envelope, increase with later type.  Consequently, it takes longer to spin down stars of later types, as observed.  Stars later than about M3, on the other hand, are fully convective.  The rotation-activity relation observed from M3 to M8 cannot, therefore, be explained by the $\\alpha\\Omega$ paradigm (the situation later than M3 is analogous to that in stars earlier than F5, except that the $\\alpha\\Omega$ dynamo fails because the star is fully convective, and not because it is mostly radiative).  In these stars, magnetic fields may also be generated by the ${\\alpha}^2$ dynamo (\\cite{Radler90}).  The ${\\alpha}^2$ dynamo is strongly dependent on Rossby number, so a rotation-activity connection may be expected.  By the argument given earlier relating spindown timescale to the depth of the convection zone, fully convective dwarfs later than M3 may also take longer to spin down than earlier M's, which is consistent with observations.  Magnetic fields in the M3 to M8 dwarfs may also be generated by a `turbulent' dynamo.  In this picture, the longer spindown timescales of M dwarfs compared to earlier types is a result of the small-scale nature of the magnetic field produced by turbulent dynamos (unlike the large-scale fields from $\\alpha\\Omega$ and ${\\alpha}^2$ dynamos), which makes angular momentum loss more difficult (\\cite{Durney}).  However, the efficiency of the turbulent dynamo is only mildly enhanced by faster rotation (\\cite{Durney}).  One might expect, therefore, a sharp weakening of the rotation-activity connection at $\\sim$ M3, where full convection sets in.  No such change is observed from M0 to M8, however (\\cite{Delfosse98}; \\cite{Mohanty02}, hereafter MB; \\cite{Mohantyprep}):  activity is seen to saturate beyond a cutoff velocity in all these objects, just as in earlier types, and the saturation level (e.g., $\\Lhal/\\Lbol \\sim$ 10$^{-4}$ for chromospheric activity) is similar to that in earlier types.  This may indicate that the ${\\alpha}^2$ dynamo actually dominates in these objects.  Alternatively, it may be that the turbulent dynamo is predominantly responsible for field generation even before the onset of full convection, so that the transition to full convection is not observationally significant.  This is not implausible given that this dynamo operates throughout the convection zone, and that even at M0, a substantial fraction of the star is convective.  However, this explanation is not fully satisfactory, given that activity down to about M6 seems to saturate at very small rotational velocities ($\\lesssim$ 3 \\kms; \\cite{Delfosse98}): it is unclear how a weakly rotation-dependent dynamo could induce saturation at such low velocities (unless it creates very strong fields to begin with).  Recently, MB have observed that the threshold velocity for saturation in \\hal surface flux ($\\Fhal$) may be somewhat higher in M6 and later dwarfs ($\\sim$ 10 \\kms), and that at lower velocities, no rotation-activity connection is apparent.  This may be attributed to the presence of a weakly rotation-dependent turbulent dynamo:  at low rotational velocities, the field generation rate, and hence activity, is effectively independent of rotation, but at high enough velocities, rotation finally becomes important, causing enhanced field generation and saturated activity.  Whether an ${\\alpha}^2$ dynamo, a turbulent one, or both are operational in these stars needs to be clarified by further observations.  In any case, whatever the details of the field generation process may be in the fully convective stars down to $\\sim$ M8, it is clear that these dwarfs have no difficulty in producing substantial activity.  As mentioned above, \\cite{Delfosse98} have shown that both saturated and unsaturated activity levels in dwarfs down to $\\sim$ M6 are similar to that in the early M dwarfs, and MB have shown that this is true even for dwarfs down to $\\sim$ M8 (with rapid rotation, i.e., \\vsini $\\gtrsim$ 10 \\kms, always leading to saturation).  This is in sharp contrast to the situation in M9 and later dwarfs, which we now briefly discuss.  \\subsection{Rotation and Activity in Late M and L Dwarfs} Gizis et al. (2000) report a turnover in the fraction of cool dwarfs with \\hal in emission.  The fraction rises from early to mid-M, so that by M7-M8, $\\sim$ 100\\% of the dwarfs are in emission.  Thereafter, though, the fraction quickly drops, to about 60\\% by L0 and 8\\% by L4.  None of their sample L5 and later shows definite signs of emission.  This is all the more remarkable, given that the bolometric flux falls sharply with decreasing \\teff.  Therefore a constant \\hal flux should be increasingly evident with later types against the fainter background.  The observations clearly imply that chromospheric activity falls off after $\\sim$ M8\\footnote{\\hal emission has been seen in a T dwarf (\\cite{Burgasser00}), which is much cooler than L dwarfs.  However, most T dwarfs do not show any emission, and this particular case is regarded as an anomaly; the emission is thought to have external causes, such as Roche lobe overflow from a companion, and not an indicator of intrinsic chromospheric activity}.  Consistent with the Gizis et al. (2000) result, MB report that activity levels (measured either in terms of $\\Fhal$ or $\\Lhal/\\Lbol$) sharply drop beginning $\\sim$ M9.  The majority of the MB L dwarfs exhibit no \\hal emission.  Moreover, even the maximum activity detected in the L's is barely equal to the minimum levels in the M5 to M8 dwarfs, and more than an order of magnitude lower than the saturated M5-M8 levels.  This is remarkable, especially since the L dwarfs, by M dwarf standards, are very rapid rotators:  almost all the L's in the MB sample have \\vsini $>$ 20 \\kms, and the fastest has \\vsini $\\approx$ 60 \\kms.  Evidently, the ``standard'' rotation-activity connection breaks down at about M9.  A variety of reasons have been proposed to explain this behavior.  Basri (2000) proposed that very fast rotation may order fluid motions, thereby damping turbulence and lowering the efficiency of any turbulent dynamo.  This no longer seems likely however, given that the L dwarf with highest \\vsini exhibits the highest chromospheric activity among the MB L dwarf sample.  Alternatively, models indicate that the convective velocities diminish as one moves from mid-M to the late M and L dwarfs (see Figs. 9 and 10).  Assuming that turbulence is driven by the convective motions, lower convective velocities may imply less energy in turbulent motions and hence a damping of the turbulent dynamo.  The latter cannot be the whole story, however, for three reasons.  One, the models suggest that convective velocities decrease continuously as one moves from mid-M to later types.  In spite of this, M6 to M8 dwarfs seem quite capable of generating substantial activity.  It is difficult to see why the diminishing convective velocity would suddenly begin to have an observable effect on the dynamo at M9 and not before (see, however, the discussion of acoustic heating in \\S 5.3).  Two, dwarfs at M9 (LP 944-20) and M9.5 (BRI 0021) that exhibit little or no basal chromospheric and coronal activity have been observed to flare (\\cite{Rutledge00}; \\cite{Reid99}).  This implies that whatever dynamo is operational in them is indeed capable of generating substantial fields, but that the energy in the field cannot usually be extracted to drive activity.  Three, as discussed above, an ${\\alpha}^2$ dynamo instead of a turbulent one may be dominant in these fully convective dwarfs.  Since the efficiency of this dynamo increases strongly with decreasing Rossby number, its effects should become stronger with decreasing \\teff (which in these objects translates to longer convective timescales) and faster rotation, not weaker.  A different solution has been proposed by Meyer \\& Meyer-Hofmeister (1999), which addresses the efficiency of converting magnetic field energy into activity, rather than the efficiency of the field-producing dynamo itself.  They suggest that the electrical resistivities in the cool, dense and predominantly neutral atmospheres of late-M and L dwarfs are very high.  Consequently, these atmospheres cannot support substantial currents.  As a result, since $\\del\\times\\B = (4\\pi/c)\\j$, large non-potential field configurations (i.e., $\\mid\\del\\times\\B\\mid \\gg 0$) are not possible in the atmosphere.  Potential fields represent the lowest energy state; without large departures from this state, the magnetic free energy available to support a chromosphere and corona and drive activity is small.  {\\it It is this scenario that we explore in the present paper}.  The above picture assumes that magnetic fields are the underlying cause of chromospheric and coronal activity in these cool dwarfs.  It is possible that other energy sources are responsible for activity instead.  The most important candidate in this regard is acoustic energy (which may, for example, play a part in supporting the lower chromosphere of the Sun).  To examine this question, we calculate acoustic fluxes as well in the \\teff range of interest here, and compare them to the observed \\hal fluxes.  In \\S 2, we describe the atmospheric conditions in mid-M to L dwarfs, and the basic relevant equations.  The latter are given in more detail in the Appendix.  In \\S 3, we derive the diffusion coefficients in these atmospheres, for a range of optical depths.  The magnitudes of the coefficients are calculated, and the dominant diffusion process found.  Using the latter, we derive the equilibrium field diffusion equation in \\S 4.  The consequences of this equation are discussed in \\S 5.  In \\S 5.1 we calculate the atmospheric Reynolds numbers.  These indicate how well the field is coupled to atmospheric motions, and the magnitude of non-potential field that can be generated by such motions.  We show in \\S 5.2 that the Reynolds numbers imply that the observed falloff in \\hal activity with \\teff may indeed be a consequence of a decline in field-atmosphere coupling.  In \\S 5.3, we discuss acoustic waves as an alternative to magnetic fields for generating activity.  Finally, in \\S 5.4 we suggest a possible scenario whereby flaring can occur in otherwise inactive ultra-cool dwarfs.  Our results are summarized in \\S 6. ",
        "conclusions": "Observations show that \\hal activity rapidly declines from mid-M to L dwarfs.  We have investigated the possibility that this results from high electrical resistivities in the cool, mostly neutral atmospheres of these objects.  Using the atmospheric models of Allard and Hauschildt, the resistivities are calculated for a range of optical depths, and effective temperatures appropriate to  $\\sim$ M5 to L6 dwarfs.  We find that the resistivities are indeed very large, predominantly due to the large rate of collisions between neutrals and charged particles.  The collisions efficiently knock the charged particles off the magnetic field lines, and effectively decouple the field from atmospheric fluid motions.  As a result, fluid motions in the atmosphere cannot generate substantial magnetic stresses.  Since large currents be sustained in this very resistive environment, magnetic stresses generated in the highly conducting interior cannot efficiently traverse the atmosphere either.  Finally, the atmospheric magnetic fields themselves may be weak if a turbulent dynamo is present: the large atmospheric resistivities would push the dynamo into the interior, and the small-scale fields it generates would not reach very far into the atmosphere.  This would further hamper the creation and transport of currents through the atmosphere.  Consequently, both the production and propagation of magnetic stresses is severely hampered in these dwarfs.  Since the atmospheric resistivities increase strongly with decreasing \\teff, these difficulties in generation and transport of currents increase as well as \\teff goes down.  Therefore the magnetic energy available to support a chromosphere and generate \\hal activity strongly diminishes as one moves from mid-M to L.  Hence, the observed decline in \\hal activity from mid-M to L may indeed be a consequence of the high atmospheric resistivities.  Our above conclusion assumes that \\hal activity is magnetically driven at spectral types M5 and later.  We cannot be certain of the validity of this assumption, since we have not calculated the actual magnetic energy flux resulting from atmospheric and interior fluid motions, and compared it to the observed \\hal fluxes.  However, in the absence of magnetic heating, acoustic waves appear to be the only viable energy source for \\hal emission.  A simple Lighthill-Proudman calculation shows, that acoustic heating is probably not sufficient to explain the peak \\hal fluxes observed in mid-M to L dwarfs, although it may still be energetically important in the unsaturated M's, and those L's in which no \\hal emission is currently detected.  Since there is a decline in the peak \\hal emission from mid-M to L dwarfs, magnetic heating probably does need to be invoked.  Even if the acoustic flux contributes to the \\hal emission, it too decreases with \\teff.  Hence it may be expected to further exaggerate the decline in emission, as one moves to cooler dwarfs, that would result from a lessening of magnetic heating alone.  We have also sketched a possible mechanism whereby flaring might occur even in the absence of continuous, long-term activity.  Our scenario involves thick, twisted flux tubes that are generated in the interior and rapidly rise up and dissipate their energy in the upper atmosphere.  It remains to be seen whether this is indeed a viable mechanism.   \\clearpage  \\appendix"
    },
    "0201/astro-ph0201204_arXiv.txt": {
        "abstract": "{ Using a new, increased dataset of 7 QSOs from VLT/UVES observations combined with one QSO from the literature, the minimum Doppler parameters as a function of neutral hydrogen column density $N_\\ion{H}{i}$, $b_\\mathrm{c}(N_\\ion{H}{i})$, of the Ly$\\alpha$ forest has been derived at three redshifts $<\\!z\\!> \\,=$ 2.1, 3.3 and 3.8. In particular, five QSOs at $<\\!z\\!>\\,= $ 2.1 enable us to study the cosmic variance of $b_\\mathrm{c}(N_\\ion{H}{i})$ at lower $z$ for the first time. When incompleteness of the number of the observed lines towards lower $N_\\ion{H}{i}$ is accounted for, the derived slopes of $b_\\mathrm{c}(N_\\ion{H}{i})$, $(\\Gamma-1)$, are consistent with no-$z$ evolution with an indication of lower value at $<\\!z\\!> \\,=$ 3.3, while $b_\\mathrm{c}(N_\\ion{H}{i})$ at a fixed column density $N_\\ion{H}{i} = 10^{13.6} \\ \\mathrm{cm}^{-2}$, $b_\\mathrm{c}(13.6)$, increases as $z$ decreases. Assuming a QSO-dominated UV background, the slope of the equation of state $(\\gamma-1)$ shows no $z$-evolution within large uncertainties and the temperature at the mean density, $T_{0}$, decreases as $z$ decreases at three redshift ranges. There is a large fluctuation of $(\\Gamma-1)$ and $b_\\mathrm{c}(13.6)$ even at the similar redshifts, in particular at $<\\!z\\!> \\,=$ 3.3 and 3.8. The lower $(\\Gamma-1)$ and higher $b_\\mathrm{c}(13.6)$ values at $z \\sim 3.1$ and 3.6  compared to ones at $z \\sim 3.4$ and 3.9 are caused by a lack of lower-$N_\\ion{H}{i}$ and lower-$b$ lines at lower-$z$ parts of each QSO at $z > 3$, probably due to the \\ion{He}{ii} reionization. This result suggests that an impact from the \\ion{He}{ii} reionization on the forest might be mainly on the lower-$N_\\ion{H}{i}$ forest. From this new dataset, we find some forest clouds with a high ratio of \\ion{Si}{iv} column density to \\ion{C}{iv} column density, $N_\\ion{Si}{iv}$/$N_\\ion{C}{iv}$, at $z < 2.5$, although the bulk of the forest clouds shows lower $N_\\ion{Si}{iv}$/$N_\\ion{C}{iv}$. This high $N_\\ion{Si}{iv}$/$N_\\ion{C}{iv}$ at $z < 2.5$ suggests that some forest clouds are exposed to a soft UV background. This lack of strong discontinuity of $N_\\ion{Si}{iv}$/$N_\\ion{C}{iv}$ at $N_\\ion{H}{i} = 10^{14-17} \\mathrm{cm}^{-2}$ at $z \\sim 3$ suggests that $N_\\ion{Si}{iv}$/$N_\\ion{C}{iv}$ might not be a good observational tool to probe the \\ion{He}{ii} reionization and/or that the UV background might be strongly affected by local, high-$z$ galaxies at $z < 3$. ",
        "introduction": "The Ly$\\alpha$ forest imprinted in the spectra of high-$z$ QSOs arises from the fluctuating low-density intergalactic medium (IGM), highly photoionized by the metagalactic UV background. Since the universe expands adiabatically and the Ly$\\alpha$ forest is in photoionization equilibrium with the UV background, the temperature of the Ly$\\alpha$ forest as a function of $z$ provides a unique and powerful tool to probe the physical state of the IGM and the reionization history of the universe (Hui \\& Gnedin \\cite{hui97}; Schaye et al. \\cite{sch99}; Ricotti, Gnedin \\& Shull \\cite{ric00}; McDonald et al. \\cite{mc00}).  For a low-density (the baryon overdensity $\\delta < \\ \\sim \\!10$), photoionized gas, the temperature of the gas is shown to be tightly correlated with the overdensity of the gas. This relation, i.e. the equation of state, is defined by $T=T_{0}(1+\\delta)^{\\gamma-1}$, where $T$ is the gas temperature in K, $T_{0}$ is the gas temperature in K at the mean gas density and $(\\gamma-1)$ is a constant at a given redshift $z$. Both $T_{0}$ and $(\\gamma-1)$ are a function of $z$, depending on the thermal history of the IGM (Hui \\& Gnedin \\cite{hui97}).  This equation of state, however, is not directly observable. Instead of $T$ and $(1+\\delta)$, observations only provide the neutral hydrogen column density $N_\\ion{H}{i}$ (in cm$^{-2}$) and the Doppler parameter $b$ (in km s$^{-1}$) of the forest absorption lines. In practice, a lower cutoff envelope in the $N_\\ion{H}{i}$--$b$ distribution is used to probe the {\\it upper}  limit on the temperature of the IGM since the forest lines could be broadened by processes other than the thermal broadening. Translating a $N_\\ion{H}{i}$--$b$ envelope into a $(1+\\delta)$--$T$ relation depends on many physical assumptions, such as the ionizing UV background $J_{\\nu}$ (Miralda-Escud\\'e et al. \\cite{mir96}; Schaye et al. \\cite{sch99}).  This minimum Doppler cutoff $b_\\mathrm{c}(N_\\ion{H}{i})$ can be described by \\begin{equation} \\label{eq2} \\log(b_\\mathrm{c}) = \\log(b_{0}) + (\\Gamma-1) \\, \\log(N_\\ion{H}{i}), \\end{equation} where $\\log(b_{0})$ is the intercept of the cutoff in the $\\log (N_\\ion{H}{i})$--$\\log b$ diagram and $(\\Gamma -1)$ is the slope of the cutoff (Schaye et al. \\cite{sch99}).   \\begin{table*} \\caption[]{Analyzed QSOs} \\label{tab1} \\begin{tabular}{lcccccl} \\hline \\noalign{\\smallskip} QSO & $z_\\mathrm{em}^{\\mathrm{a}}$ & mag$^{\\mathrm{a}}$ & $\\lambda\\lambda$ (\\AA\\/) & $z_\\mathrm{Ly\\alpha}$ & \\# of lines$^{\\mathrm{b}}$ & Comments \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\object{Q1101--264} & 2.145 & 16.0 & 3500--3778 & 1.88--2.11 & 69 & UVES SV, a damped system at $z=1.8386$ \\\\ \\object{J2233--606} & 2.238 & 17.5 & 3500--3890 & 1.88--2.20 & 88 & UVES Commissioning I\\\\ \\object{HE1122--1648} & 2.400 & 17.7 & 3500--4091 & 1.88--2.37 & 179 & UVES SV, split into 2$^{\\mathrm c}$\\\\ \\object{HE2217--2818} & 2.413 & 16.0 & 3510--4100 & 1.89--2.37 & 159 & UVES Commissioning I, split into 2$^{\\mathrm d}$\\\\ \\object{HE1347--2457} & 2.534 & 16.8 & 3760--4100 & 2.09--2.37 & 91 & UVES SV, incomplete observations \\\\ \\object{Q0302--003} & 3.281 & 18.4 & 4808--5150 & 2.96--3.24 & 107 & UVES Commissioning I, incomplete observations \\\\ \\object{Q0055--269} & 3.655 & 17.9 & 4850--5598 & 2.99--3.60 & 264 & UVES, Sept 20--22, 2000, split into 2$^{\\mathrm e}$\\\\ \\object{Q0000--263} & 4.127 & 17.9 & 5450--6100 & 3.48--4.02 & 209 & Lu et al. (\\cite{lu96}), split into 2$^{\\mathrm f}$ \\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\begin{list}{}{} \\item[$^{\\mathrm{a}}$] Taken from the SIMBAD database. The magnitude of \\object{HE1347--2457} is from NED. \\item[$^{\\mathrm{b}}$] For $N_\\ion{H}{i} = 10^{12.5-14.5} \\mathrm{cm}^{-2}$. Only for lines with the errors less than 25\\% in both $N_\\ion{H}{i}$ and $b$. \\item[$^{\\mathrm{c}}$] For Sample B, the spectrum is split into 3500--3800 \\AA\\/ (87 lines) and 3800--4091 \\AA\\/ (92 lines). \\item[$^{\\mathrm{d}}$] For Sample B, the spectrum is split into 3510--3800 \\AA\\/ (76 lines) and 3800--4100 \\AA\\/ (83 lines). \\item[$^{\\mathrm{e}}$] For Sample B, the spectrum is split into 4850--5220 \\AA\\/ (110 lines) and 5220--5598 \\AA\\/ (154 lines). \\item[$^{\\mathrm{f}}$] For Sample B, the spectrum is split into 5450--5820 \\AA\\/ (106 lines) and 5820--6100 \\AA\\/ (103 lines). \\end{list} \\end{table*}  From observations alone, both no $z$-evolution of $N_\\ion{H}{i}$-independent $b_\\mathrm{c}$ (Kirkman \\& Tytler \\cite{kir97}; Savaglio et al. \\cite{sav99}) and increasing $b_\\mathrm{c}$ with decreasing $z$ (Kim et al. \\cite{kim97}) have been claimed. Results from simulations combined with observations have also claimed both no-$z$ evolution of $T_{0}$ and $(\\gamma-1)$ (McDonald et al. \\cite{mc00}) and a $z$-evolution (Ricotti et al. \\cite{ric00}; Schaye et al. \\cite{sch00}; Kim, Cristiani \\& D'Odorico \\cite{kim01a}). Deriving $b_\\mathrm{c}(N_\\ion{H}{i})$ from observations depends on many factors such as the method of line deblending, the number of available absorption lines, the metal-line contamination, and the method of fitting the lower $N_\\ion{H}{i}$--$b$ envelope (Hu et al. \\cite{hu95}; Kirkman \\& Tytler \\cite{kir97}; Bryan \\& Machacek \\cite{bry00}; McDonald et al. \\cite{mc00}; Ricotti et al. \\cite{ric00}; Shaye et al. \\cite{sch00}; Kim et al. \\cite{kim01a}). The different approaches and the limited numbers of lines have led, in part, to the contradicting results on the evolution of $b_\\mathrm{c}(N_\\ion{H}{i})$ in the literature.  Here, using a new, increased dataset from 7 QSOs observed with the VLT/UVES combined with the published data on one QSO obtained with Keck/HIRES, we present the evolution of the Doppler cutoff $b_\\mathrm{c}(N_\\ion{H}{i})$ at three redshifts $<\\!z\\!>\\,= $ 2.1, 3.3 and 3.8. In particular, five QSOs at $<\\!z\\!>\\,= $ 2.1 enable us to study the cosmic variance of $b_\\mathrm{c}(N_\\ion{H}{i})$ and to improve a determination of $b_\\mathrm{c}(N_\\ion{H}{i})$ at lower $z$ for the first time. In Sect. 2, we briefly describe the data used in this study. The analyses of the observations are presented in Sect. 3. The discussion is in Sect. 4 and the conclusions are summarized in Sect. 5. In this study, all the quoted uncertainties are $1\\sigma$ errors.   \\begin{table*} \\caption[]{The power-law fits to $b_\\mathrm{c}(N_\\ion{H}{i})$} \\label{tab2} \\begin{tabular}{ccccccccc} \\hline \\noalign{\\smallskip} \\multicolumn{9}{c}{Sample A}\\\\ \\hline \\noalign{\\smallskip} $<\\!z\\!>$ & $\\log N_\\ion{H}{i}$ & \\# of lines & $\\log(b_\\mathrm{0,i})$ & $(\\Gamma-1)_\\mathrm{i}$ & $b_\\mathrm{c,i,}(13.6)$ & $\\log(b_\\mathrm{0,s})$ & $(\\Gamma-1)_\\mathrm{s}$ & $b_\\mathrm{c,s}(13.6)$ \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} 2.1 & 13.0--14.5 & 349 & $-0.745 \\pm 0.089$ & $0.150 \\pm 0.006$ & $19.8 \\pm 0.8$ & $-0.495 \\pm 0.074$ & $0.131 \\pm 0.002$ & $19.1 \\pm 1.1$ \\\\ 3.3 & 13.0--14.5 & 275 & $-0.413 \\pm 0.116$ & $0.122 \\pm 0.008$ & $17.2 \\pm 1.0$ & $0.279 \\pm 0.473$ & $0.072 \\pm 0.127$ & $18.4 \\pm 1.5$ \\\\ 3.8 & 13.3--14.5 & 152 & $-0.948 \\pm 0.125$ & $0.159 \\pm 0.008$ & $16.8 \\pm 1.6$ & $-2.699 \\pm 0.159$ & $0.285 \\pm 0.042$ & $14.9 \\pm 1.1$ \\\\ 2.0$^{\\mathrm a}$ & 13.0--14.5 & 176 & $-0.214 \\pm 0.180$ & $0.111 \\pm 0.013$ & $19.1 \\pm 1.1$ & $-1.078 \\pm 0.173$ & $0.172 \\pm 0.047$ & $18.0 \\pm 1.1$ \\\\ 2.2$^{\\mathrm a}$ & 13.0--14.5 & 173 & $-0.954 \\pm 0.132$ & $0.167 \\pm 0.010$ & $19.7 \\pm 1.0$ & $0.400 \\pm 0.110$ & $0.066 \\pm 0.030$ & $19.9 \\pm 1.1$ \\\\ 3.1$^{\\mathrm b}$ & 13.0--14.5 & 157 & $0.645 \\pm 0.379$ & $0.048 \\pm 0.027$ & $19.9 \\pm 1.5$ & $-0.366 \\pm 0.103$ & $0.122 \\pm 0.028$ & $19.6 \\pm 1.1$ \\\\ 3.4$^{\\mathrm b}$ & 13.0--14.5 & 118 & $-0.330 \\pm 0.053$ & $0.117 \\pm 0.009$ & $17.5 \\pm 0.8$ & $-0.745 \\pm 0.131$ & $0.142 \\pm 0.035$ & $15.5 \\pm 1.1$ \\\\ 3.6$^{\\mathrm c}$ & 13.3--14.5 & 74 & $0.572 \\pm 0.331$ & $0.054 \\pm 0.023$ & $20.4 \\pm 0.8$ & $1.092 \\pm 0.007$ & $0.014 \\pm 0.002$ & $19.0 \\pm 1.0$ \\\\ 3.9$^{\\mathrm c}$ & 13.3--14.5 & 78 & $-0.697 \\pm 0.133$ & $0.143 \\pm 0.010$ & $17.5 \\pm 0.7$ & $-0.499 \\pm 0.235$ & $0.122 \\pm 0.062$ & $14.5 \\pm 1.1$ \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\multicolumn{9}{c}{Sample B (averaged for the individual members)} \\\\ \\hline \\noalign{\\smallskip} 2.1 & 13.0--14.5 &  ...  & $-0.044 \\pm 0.506$ & $0.102 \\pm 0.037$ & $21.4 \\pm 0.8$ & $0.053 \\pm 0.731$ & $0.090 \\pm 0.053$ & $19.2 \\pm 1.1$ \\\\  3.3 & 13.0--14.5 &  ...  & $0.617 \\pm 0.966$ & $0.051 \\pm 0.067$ & $20.3 \\pm 2.7$ & $0.283 \\pm 0.970$ & $0.072 \\pm 0.068$ & $18.3 \\pm 2.7$ \\\\  3.8 & 13.3--14.5 &  ...  & $-0.063 \\pm 0.898$ & $0.098 \\pm 0.062$ & $19.0 \\pm 2.1$ & $0.182 \\pm 1.350$ & $0.078 \\pm 0.094$ & $17.4 \\pm 3.1$ \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\multicolumn{9}{c}{Results from Schaye et al. (\\cite{sch00}) (their sample corresponds to our Sample B)}\\\\ \\hline \\noalign{\\smallskip} $\\sim 3.1$ & 12.5--14.5 & ... & ... & $\\le 0.150$ & $\\sim 23$--24 & ... & ... & ... \\\\ $\\sim 3.1$ & 12.5--14.5 & ... & ... & $\\sim 0$ & $\\sim 20$--22 & ... & ... & ... \\\\ $\\sim 3.8$ & 12.5--14.8 & ... & ... & $\\le 0.150$ & $\\sim 18$--19 & ... & ... & ... \\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\begin{list}{}{} \\item[$^{\\mathrm{a}}$] In order to have a similar number of lines at $z \\sim 2.0$ and $z \\sim 2.2$, the $<\\!z\\!> \\ = 2$ sample consists of the line parameters from \\object{Q0011--264}, \\object{J2233--606}, \\object{HE1122--1648} at 3500--3760 \\AA\\/ and \\object{HE2217--2818} at 3510--3760 \\AA\\/. The the $<\\!z\\!> \\ = 2.2$ sample consists of the line parameters from \\object{Q1347--2457}, \\object{HE1122--1648} at 3760--4100 \\AA\\/ and \\object{HE2217--2818} at 3760--4100 \\AA\\/. \\item[$^{\\mathrm{b}}$] The $<\\!z\\!> \\ = 3.1$ sample and the $<\\!z\\!> \\ = 3.4$ sample are taken from \\object{Q0302--003} and \\object{Q0055--269} at 4850--5220 \\AA\\/, from \\object{Q0055--269} at 5220--5598 \\AA\\/, respectively. \\item[$^{\\mathrm{c}}$] The $<\\!z\\!> \\ = 3.6$ sample and the $<\\!z\\!> \\ = 3.9$ sample are taken from 5450--5820 \\AA\\/ and from 5820--6100 \\AA\\/, respectively. \\end{list} \\end{table*} ",
        "conclusions": "Using the new, large dataset from high S/N, high resolution VLT/UVES data combined with one Keck/HIRES QSO in the literature, the minimum cutoff Doppler parameter as a function of $N_\\ion{H}{i}$, $b_\\mathrm{c}(N_\\ion{H}{i})$, of the Ly$\\alpha$ forest has been derived at $<\\!z\\!> \\,=$ 2.1, 3.3 and 3.8. We have found:  \\begin{enumerate}  \\item When incompleteness is accounted for, the derived $(\\Gamma-1)$ is consistent with no-$z$ evolution with a suggestion of lower $(\\Gamma-1)$ at $z \\sim 3.1$, while $b_\\mathrm{c}(13.6)$ increases as $z$ decreases. These results suggest that $(\\gamma-1)$ shows no $z$-evolution within uncertainties and that $T_{0}$ decreases as $z$ decreases, assuming the QSO-dominated UV background from Haardt \\& Madau (\\cite{haa96}).  \\item The $(\\Gamma-1)$ and $b_\\mathrm{c}(13.6)$ values show a large fluctuation when derived using a subsample even at the similar redshifts. The fluctuation is larger at $z > 3$ than at $z < 3$. Although smaller number of lines in the fit could result in this fluctuation, we use a similar number of lines to derive $(\\Gamma-1)$ and $b_\\mathrm{c}(13.6)$ for each subsample. This result suggests that there might be a large fluctuation in the IGM temperature along different sightlines even at similar $z$.  \\item There is a suggestion of a flatter $(\\Gamma-1)$ and higher $b_\\mathrm{c}(13.6)$ along each line of sight as $z$ decreases at $z > 3$ more than 3$\\sigma$, probably due to the \\ion{He}{ii} reionization. At $z \\sim 2.1$, there is no such a significant trend along each line of sight. This result occurs due to the lack of lower-$N_\\ion{H}{i}$ and smaller-$b$ lines in the lower-$z$ part of the spectra at $z > 3$. Our result implies that the impact of the \\ion{He}{ii} reionization on the Ly$\\alpha$ forest might be mainly on the lower-$N_\\ion{H}{i}$ forest and that its significance might be smaller than previously suggested.  \\item The lack of strong discontinuity of $N_\\ion{Si}{iv}$/$N_\\ion{C}{iv}$ at $N_\\ion{H}{i} = 10^{14-17} \\mathrm{cm}^{-2}$ at $z \\sim 3$ suggests that $N_\\ion{Si}{iv}$/$N_\\ion{C}{iv}$ might not be a good observational tool to probe the \\ion{He}{ii} reionization and that the UV background might be strongly contributed by local, high-$z$ galaxies at $z < 3$. \\end{enumerate}"
    },
    "0201/astro-ph0201048_arXiv.txt": {
        "abstract": "X-ray transience is the most extreme form of variability observed in AGN or normal non-active galaxies. While factors of 2-3 on timescales of days to years are quite common among AGN, X-ray transients appear only once and vanish from the X-ray sky years later. The ROSAT All-Sky Survey with its sensitivity to energies down to 0.1 keV was the an ideal tool to discover these sources. X-ray transience in AGN or galaxies can be caused by dramatic changes in the accretion rate of the central black hole or by changes of the properties of the accretion disk. So far only a handful of sources are known. In order to estimate how often such an event occurs in a galaxy, a new soft X-ray survey is needed. In these proceedings I describe the currently known X-ray transient AGN and galaxies and will argue for a new soft X-ray survey in order to discover more of these extreme X-ray sources. ",
        "introduction": "The ROSAT All-Sky Survey (RASS, \\cite{dgrupe-WB1:vog99}) has established a new phenomenon among AGN and galaxies - X-ray transience. These X-ray transient AGN and galaxies are very bright once and appear to be fainter or even vanish in X-rays when observed years later(e.g. \\cite{dgrupe-WB1:gr01a}, \\cite{dgrupe-WB1:ko99a}). X-ray transience can appear in Active Galactic Nuclei as well as in normal non-active galaxies. While the X-ray transience in AGN might be caused by changes of the accretion disk properties, in non-active galaxies it might be caused by an X-ray outburst. These X-ray outbursts can be caused by an accretion event, either by instabilities in the accretion disk or by the disruption of a star by the central black hole. Disk instabilities may cause an X-ray outburst in AGN while the tidal disruption model is favoured to explain the outbursts in non-active galaxies which do not show any signs of nuclear activity.  So far only in one source a response in the optical emission lines has been observed, IC 3599. Future soft X-ray surveys will allow us to search for this type of sources. Fast follow-up observations in the optical and in X-rays will provide us with a powerful tool to map the inner region of an AGN in order to locate the line emitting regions and to describe what the geometry of the whole system is. ",
        "conclusions": "The ROSAT All-Sky Survey has shown the potential of soft X-ray surveys to discover X-ray transient AGN and galaxies. There are in principle two types of transience in AGN and galaxies: a) a sudden decrease of the X-ray flux in normally bright X-ray sources, and b) an X-ray outburst caused by either an accretion disk instability or by a tidal disruption of a star orbiting the central black hole. The first case, which might have been seen in WPVS 007 and RX J2217.9--5941, might be the transition between a high into a low state. Such behaviour is known from galactic black hole candidates. The second case, an X-ray outburst, was observed in IC 3599, NGC 5905, RX J1624.9+7554, RX J1242.6--1119, and RX J1420.4+5334. The follow-up optical observations of IC 3599 showed how the light front moved through the inner region of the AGN/galaxy and caused different emission lines to show up over time. While reverberation mapping is a powerful tool in bright, `normal' AGN to map the inner region, X-ray transient AGN and galaxies show much larger and dramatic changes in their optical spectra. The only problem is, such X-ray outburst events are rare. In order to find those extraordinary sources, soft X-ray surveys are needed.  New soft X-ray surveys will provide us with a statistically relevant data set that can clarify how common the chance between low and high states is in AGN. Repeating soft X-ray surveys will also give us an estimate how often outburst events appear in non-active galaxies, or in other words, how many years we have to wait until such an event happens in a galaxy. This result is also of interest for our own galaxy, which black hole mass of 2.6 10$^6$\\msun is comparable to the ones found in the galaxies in which an outburst occurred.  The requirements for future soft X-ray surveys are being sensitive at energies $<$0.3 keV and using an imaging X-ray telescope with a high effective area and a large field of view. A sensitivity in soft X-rays ($<$0.3 keV, at least) is needed, because X-ray outbursts have a very soft X-ray spectrum, due to the high accretion rate during the event. A high effective area is needed for getting enough photons to derive spectra and lightcurves on short timescales. And last but not least, a large field of view is needed to get a long coverage of the source when passing through the field of view of the telescope.  Currently, three X-ray surveys are planned, ROSITA, Lobster Eye, and MAXI. They all have in common to work at higher energies. While ROSITA  may have enough spatial resolution to identify the optical counterpart of an X-ray transient, Lobster Eye is an all-sky monitor which does not have enough spatial resolving power. MAXI will not have enough spatial resolution either, plus will not be sensitive to low energy photons. ROSITA will be a powerful tool to discover obscured hard AGN, but inefficient for the search of short living soft X-ray events such as X-ray outburst. The other problem is that the number of photons collected per scan is too small to derive lightcurves and spectra. Therefore, none of the currently planned X-ray survey missions fullfill the requirements for a search for soft X-ray transient AGN and galaxies. Therefore a new soft X-ray survey mission is needed, performed in a similar way as the RASS was, except it should be repeated several times in order to a) get long-term light curves, b) see the same number of `turn-ons' as `turn-offs'. Such a survey will not be only important for AGN/galaxy research, it is also needed for the discovery and study of e.g. new cataclysmic variables and super-soft X-ray sources."
    },
    "0201/astro-ph0201081_arXiv.txt": {
        "abstract": "New virial black-hole mass estimates are presented for a sample of 72 AGN covering three decades in optical luminosity. Using a model in which the AGN broad~-~line region (BLR) has a flattened geometry, we investigate the $M_{bh}-$~$L_{bulge}$ relation for a combined 90-object sample, consisting of the AGN plus a sample of 18 nearby inactive elliptical galaxies with dynamical black-hole mass measurements. It is found that, for all reasonable mass-to-light ratios, the $M_{bh}-L_{bulge}$ relation is equivalent to a linear scaling between bulge and black-hole mass. The best-fitting normalization of the $M_{bh}-M_{bulge}$ relation is found to be $M_{bh}=0.0012M_{bulge}$, in agreement with recent black-hole mass studies based on stellar velocity dispersions. Furthermore, the scatter around the $M_{bh}-L_{bulge}$ relation for the full sample is found to be significantly smaller than has been previously reported ($\\Delta \\log M_{bh}=0.39$ dex). Finally, using the nearby inactive elliptical galaxy sample alone, it is shown that the scatter in the $M_{bh}-L_{bulge}$ relation is only $0.33$ dex, comparable to that of the $M_{bh}-\\sigma$ relation. These results indicate that reliable black-hole mass estimates can be obtained for high redshift galaxies. ",
        "introduction": "\\label{intro} The correlation between black-hole mass and bulge luminosity is now well established for both active and inactive galaxies (Magorrian et al. 1998; Laor 1998). However, despite recent attention in the literature, the usefulness of the $M_{bh}-L_{bulge}$ relation as a black-hole mass estimator is at present severely limited due to its large scatter ($\\simeq0.5$ dex). Although the correlation between black-hole mass and stellar velocity dispersion for nearby inactive galaxies displays a much smaller scatter ($\\simeq0.3$ dex, Merritt \\& Ferrarese 2001a), it is clear that a $M_{bh}-L_{bulge}$ correlation with reduced scatter would be highly desirable, given the extreme difficulty in obtaining stellar velocity dispersions for high redshift galaxies.  This conference proceeding presents the main results of a new study (McLure \\& Dunlop 2002) in which we investigate the black hole - bulge mass relation using a 90-object sample comprised of 72 AGN (53 QSOs and 19 Seyfert 1s) and 18 nearby quiescent ellipticals with dynamically determined black-hole mass estimates. Those interested in the details of our analysis, particularly the flattened geometry model adopted for the calculation of the virial black-hole mass estimates, are referred to McLure \\& Dunlop (2002). ",
        "conclusions": "The main conclusions of this study can be summarized as follows: \\begin{itemize} \\item{The best-fitting  $M_{bh}-L_{bulge}$ relation to the combined sample of 72 AGN and 18 nearby inactive elliptical galaxies is found to be consistent with a linear scaling between black hole and bulge mass ($M_{bh}\\propto M_{bulge}^{0.95\\pm0.05}$), and to have much lower scatter than previously reported ($\\Delta \\log M_{bh}=0.39$ dex).} \\item{The best-fitting normalization of the $M_{bh}-M_{bulge}$ relation is found to be $M_{bh}=0.0012M_{bulge}$, in excellent agreement with recent stellar velocity dispersion studies.} \\item{In contrast to previous reports it is found that the scatter around the $M_{bh}-L_{bulge}$ and $M_{bh}-\\sigma$ relations for nearby inactive elliptical galaxies are comparable, at only $\\sim 0.3$ dex. } \\end{itemize}"
    },
    "0201/astro-ph0201424_arXiv.txt": {
        "abstract": "Submillimeter (and in some cases millimeter) wavelength continuum measurements are presented for a sample of 40 active galactic nuclei (probably all quasars) lensed by foreground galaxies.  The object of this study is to use the lensing boost, anywhere from $\\sim 3- 20$ times, to detect dust emission from more typical AGNs than the extremely luminous ones currently accessible without lensing.  The sources are a mix of radio loud and radio quiet quasars, and, after correction for synchrotron radation (in the few cases where necessary), 23 of the 40 (58\\%) are detected in dust emission at 850\\mic ; 11 are also detected at 450\\mic . Dust luminosities and masses are derived after correction for lensing magnification, and luminosities are plotted against redshift from $z = 1$ to $z = 4.4$, the redshift range of the sample.  The main conclusions are (1) Monochromatic submillimeter luminosities  of quasars are, on average, only a few times greater than those of local IRAS galaxies; (2) Radio quiet and radio loud quasars do not differ significantly in their dust lumimosity; (3) Mean dust luminosities of quasars and radio galaxies over the same redshift range are comparable; (4) Quasars and radio galaxies alike show evidence for more luminous and massive dust sources toward higher redshift, consistent with an early epoch of formation and possibly indicating that the percentage of obscured AGNs increases with redshift. ",
        "introduction": "Strong dust emission at infrared and submillimeter wavelengths is now recognized as a fundamental tracer of activity in galaxies, be it star formation or processes related to nuclear black holes.  Ordinary, quiescent galaxies like the Milky Way exhibit dust luminosities of order $10^{10} L_{\\sun}$, but active systems, such as ultraluminous infrared galaxies (ULIRGs), quasars, and radio galaxies can be as much as $10^3$ times more powerful (Sanders \\& Mirabel 1996).  In these objects the intense dust emission appears to be triggered by galaxy interactions and mergers, which are thought to be responsible at least in part for the formation and growth of massive elliptical galaxies and the supermassive black holes found in their nuclei (e.g., Kormendy \\& Sanders 1992; van Dokkum et al.\\ 1999).  In the mid-1980s the IRAS mission discovered many infrared-luminous galaxies and quasars at low redshifts in the mid- to far-infrared (12--100\\mic ), establishing dust emission as an important if not dominant component of the bolometric luminosity in certain classes of objects.  The dust luminosity is a result of reprocessing of shorter wavelength optical/UV radiation originating either from hot young stars, in the case of starbursts, or from the release of gravitational energy as material accretes onto a nuclear black hole, in the case of quasars.  In some galaxies both starburst and quasar activity are likely to contribute, and there remains some dispute over which process dominates in various types of objects.  Although the findings of IRAS were revolutionary for the local universe, its sensitivity was such that only a few high-redshift systems were detected: the quasar PG 1634+607 ($z = 1.3$), the ULIRG/buried quasar IRAS F10214+4724 ($z = 2.3$), the Cloverleaf quasar (H 1413+117; $z = 2.6$), and the $z = 3.9$ broad absorption line quasar APM 08279+5255 (Rowan-Robinson et al.\\ 1991; Barvainis et al.\\ 1995; Lewis et al.\\ 1998). The latter three are strongly magnified by gravitational lensing, boosting their fluxes above the IRAS detection threshholds at 60 and 100\\mic . Further progress was made at submillimeter wavelengths at the James Clerk Maxwell Telescope (JCMT), first using the single-channel photometer UKT14, then later with the submillimeter camera SCUBA.  F10214+4724 and the Cloverleaf were detected at 800 and 450\\mic\\ (Rowan-Robinson et al.\\ 1993; Barvainis et al.\\ 1992), to be followed by several other quasars and radio galaxies (Hughes et al.\\ 1997).  To date the census of published submillimeter detections of high redshift objects numbers $\\sim 150$ (mostly unidentified), some of which are either strongly lensed by individual galaxies or weakly lensed by galaxy clusters (Smail, Ivison, \\& Blain 1997).  \\begin{figure*} \\figurenum{1} \\epsscale{1.2} \\plotone{fig1.eps} \\caption{Modeled magnification versus source diameter for some of the lensed quasars in the sample.  The model for the intervening lens, based on optical data, is applied to successively larger sources at the quasar to determine the net magnification.  Magnifications that are high for the optical emission (essentially a point source) tend to drop to about 20 as the source size increases to $\\sim 200$ pc (a likely minimum dust source size).  At low magnifications the magnification tends to be constant with source size (unlabelled examples plotted include SBS 0909+532, BRI 0951+2635, LBQS 1009$-$0252, HST 12531$-$2914, HST 14113+5211, H 1413+117, and HST 14176+5226).  Model results generated by B.  McLeod (private communication). } \\end{figure*}  The high sensitivity of SCUBA, combined with the flux boost provided by strong lensing, presents an opportunity to delve deeper into the submillimeter luminosity function than would otherwise be possible.  The majority of unlensed quasars are undetectable with current instruments, so up to now {\\it typical} dust luminosties and masses of active galaxies at high redshift have remained unknown. Only the very highest luminosities can be detected without lensing.  However, there are now about 50 known strongly lensed objects, most of which are in fact quasars (see the CASTLeS compilation, http://cfa-www.harvard.edu/castles).  This paper reports the results of a submillimeter survey of a large fraction of the known lensed quasars, presenting data on a sample of 40 objects in which dust emission has been clearly detected at 850 and/or 450\\mic , or meaningfully constrained by upper limits.  These lensed quasars, while not a complete sample by any definitive criteria, are nevertheless a random, passively selected sample of quasars at high redshift -- they are quasars that happen to lie behind massive galaxies.  With a typical lensing boost of an order of magnitude, they provide a unique window on the dust content of active galaxies in the early universe. ",
        "conclusions": "The present submillimeter survey of lensed quasars shows that the dust luminosities/masses of quasars are comparable to those of radio galaxies (Archibald et al.\\ 2001) in the same redshift range ($1 \\lesssim z \\lesssim 4.4$), and that quasars, like radio galaxies, appear to have increasingly powerful dust sources toward higher redshifts.  The mean monochromatic 850\\mic\\ luminosity of high-redshift quasars is only a few times higher than that of a local sample of IRAS-selected galaxies, but because quasars have more power in the mid-infrared their total infrared luminosities ($8-1000$\\mic ) are substantially greater.  The concordance of mean dust luminosity and luminosity evolution with redshift between quasars and radio galaxies is broadly consistent with, and supportive of, obscuration/orientation unified schemes wherein these classes are by and large similar except for chance viewing angle.  The submillimeter dust flux is likely to be optically thin and therefore orientation independent.  The presence or absence of a strong radio source appears to have little effect on the submillimeter power once synchroton contamination has been accounted for, and this argues against large differences in the total masses of radio quiet versus radio loud systems.  The tentative evidence presented here for higher dust luminosities and inferred dust masses at high redshift could have several interpretations.  One possibility is that the increase is illusory, for example if dust properties change over cosmic time in such a way as to cause an increased emissivity per unit dust mass at high $z$, or if the spectral energy distribution of grain emission changes with time causing the {\\it derived} relative luminosities/masses to change. The template SED used here is not necessarily the correct one in general, and cannot be correct for all specific objects since galaxy studies have shown a fair amount of variation in the far-IR/submillimeter SEDs from one object to the next (see, e.g., Figure 2 of Adelberger \\& Steidel 2000).  Other SED shapes can dampen the trend with redshift (while yet others can strengthen it). Greenberg \\& Shen (2000) have stressed that dust in the early universe could be much different than it is today, but the form of any cosmic evolution of mean dust properties from $z=4$ to $z=1$ is anybody's guess.  Dust sources in quasars are powered by absorption of energy from the central AGN and by star formation, with the proportionate mix between the two being in general unknown. An increased mass of dust in early galaxies could come about because of more frequent galaxy interactions and mergers when the universe was smaller and more dense. These mergers could, in addition to dumping mass into the central regions of the host galaxy, trigger vigorous star formation to power the dust.  Alternatively, if the primary energy source for the grains is the central AGN, the increased luminosity could be caused by a larger dust covering factor at higher $z$, meaning that more of the quasar light would be intercepted and reprocessed into IR/submillimeter emission.  A larger covering factor could result from a larger total dust mass, or simply a larger fraction of the dust being exposed to the primary radiation.  If there is a torus surrounding the AGN, it could be physically thicker (i.e., smaller opening angle).  A second possibility is that material has not had a chance to settle into an optically thick torus because of the youth of the systems and their more violent environments, leaving relatively more of the grains exposed to the nuclear radiation.  The larger covering factor interpretation imples a larger fraction of obscured AGNs at high redshifts, as a lower proportion of direct quasar light escapes to distant observers.  Models for the x-ray background are still uncertain, but recent work by Gilli et al.\\ (2001) finds a better fit to the currently available x-ray data with an increasing ratio of obscured to unobscured AGNs toward higher redshifts.  That particular model seems to indicate that the evolution stops above about $z = 1.5$. The increasing dust luminosity with redshift found by Archibald et al.\\ (2001) for radio galaxies, and the similar relation found here for quasars, suggests that the proportion of obscured AGNs may continue to increase out to $z = 2$ or beyond."
    },
    "0201/astro-ph0201338_arXiv.txt": {
        "abstract": "Early-type galaxies exhibit thermal and molecular resonance  emission from dust that is shed and heated through stellar mass loss as a subset of the population moves through the AGB phase of evolution. Because this emission can give direct insight into stellar evolution in addition to galactic stellar mass loss and ISM injection rates, we conducted a program to search for this signature emission with CAM on ISO. We obtained 6-15\\( \\mu  \\)m imaging observations in six narrow bands for nine elliptical galaxies; every galaxy is detected in every band. For wavelengths shorter than 9\\( \\mu  \\)m, the spectra are well matched by a blackbody, originating from the K and M stars that dominate the integrated light of elliptical galaxies. However, at wavelengths between 9\\( \\mu  \\)m and 15\\( \\mu  \\)m, the galaxies display excess emission relative to the stellar photospheric radiation. Additional data taken with the fine resolution circular variable filter on one source clearly shows broad emission from 9\\( \\mu  \\)m to 15\\( \\mu  \\)m, peaking around 10\\( \\mu  \\)m. This result is consistent with the known, broad silicate feature at 9.7\\( \\mu  \\)m, originating in the circumstellar envelopes of AGB stars. This emission is compared with studies of Galactic and LMC AGB stars to derive cumulative mass loss rates. In general, these mass loss rates agree with the expected \\( \\sim  \\)0.8 M\\( _{\\sun }yr^{-1} \\) value predicted by stellar evolutionary models. Both the photospheric and circumstellar envelope emission follow a de Vaucouleurs' R\\( ^{1/4} \\) law, supporting the conclusion that the mid-infrared excess emission originates in the stellar component of the galaxies and acts as a tracer of AGB mass loss and mass injection into the ISM. ",
        "introduction": "Elliptical galaxies are comprised of old stellar populations in which low mass stars evolve off the main sequence, eventually becoming white dwarfs. These post-main sequence stars experience an asymptotic giant branch (AGB) phase where up to 0.3 M\\( _{\\sun } \\) is shed into the Interstellar Medium (ISM). As this mass is ejected from the star, dust condenses out of the gas in a cool stellar wind and forms an envelope at a few to ten stellar radii. (For a review of AGB stars and circumstellar envelopes see \\citet{AGB Star Review}). Dust acts as an agent to this mass loss as it transfers radiation momentum from the star to the gas through collisions. The stellar radiation absorbed by the dust heats it to temperatures from 300-1000K, depending on the distance from the star. Subsequently, the dust cools thermally and has been detected and explored through the Infrared Astronomical Satellite (IRAS) and the Infrared Space Observatory (ISO) wide band filters at 12, 25, 60 and 100\\( \\mu  \\)m. (See \\citet{Jura} for an early study of elliptical galaxies with IRAS observations. See \\citet{MidIR View of Galaxies} for a recent review of mid and far IR emission from all types of galaxies.) In addition to this thermal emission, broad line absorption or emission is observed in most AGB stars at 10-12\\( \\mu  \\)m and for oxygen-rich AGB stars, again at 20\\( \\mu  \\)m (See \\citet{Speck} for a recent study of AGB mid-IR features). For older, low mass stars found in elliptical galaxies, the AGB star envelopes have a high oxygen content as opposed to the high carbon content present in high mass AGB envelopes, which occurs through inner layer carbon dredge up and expulsion. The theoretical models of dust condensation show amorphous silicates (Si-O) to be the major constituent of the dust in these oxygen-rich environments, where the stretching and bending of the Si-O bond is responsible for the 10 and 20\\( \\mu  \\)m absorption/emission. (Hereafter AGB, refers to the low mass, oxygen-rich subclass of AGB stars.) One of the surprises revealed by ISO are that in addition to amorphous silicates, crystalline silicates are observed through many narrow band features from 15-45\\( \\mu  \\)m \\citep{Ref A}. These observations have sparked renewed interest in dust particle formation in AGB environments \\citep{Ref B,Ref C}. The dust composition, size distribution and other properties (albedo, dipole strength, etc) define the interaction with the gas and ultimately control the observational consequences. Consequently, much laboratory work is being devoted to ``growing'' silicate dust particles and attempting to extract similar particles from interplanetary dust particles brought to Earth through comets and meteorites \\citep{Ref D,Ref E}. These theoretical and lab studies have led to a working model of AGB star dust formation and emission.   Although much work has been done on the study of individual AGB stars, population studies of these AGB dust features in early-type galaxies is relatively unexplored due to the faint nature of the emission involved. In an elliptical galaxy, it is the sum of many individual oxygen-rich AGB stars that will produce in aggregate a similar feature to individual AGB stars. This feature can be used to confirm this general stellar evolution picture and determine the mass loss rate into the entire galaxy. An effort to detect this mid-IR excess in early-type galaxies was first carried out by \\citet{KGWW} (KGW hereafter) who used the IRAS All Sky Survey coupled with ground based 10\\( \\mu  \\)m data to search for the dusty component of the ISM in nearby elliptical galaxies. Because this emission is quite faint, these galaxies were detected at low signal-to-noise ratio. Nevertheless, they observed excess emission at 12\\( \\mu  \\)m relative to a derived stellar continuum, indicating emission from the circumstellar dust envelope. KGW scaled the emission from Galactic AGB stars to their signal and they determined galaxy-wide mass loss rates of $\\sim$0.7M\\( _{\\sun }yr^{-1} \\). Also, they show that the dust emission is extended on the scale of the galaxy, ruling out the hypothesis that the dust emission originates only in the nuclear regions. However, many of the galaxy detections were at the 2\\( \\sigma  \\) level and the coarse IRAS sampling (1\\( \\arcmin  \\)x5\\( \\arcmin  \\)) prevented good spatial sampling of the galaxies. Also in the KGW study, the sample of Galactic stars used to calibrate the observed excess did not have well determined distances, leading to errors in the scaling relation. Subsequent to the KGW work, significant advances in understanding have been made in the study of individual AGB stars with detailed observational monitoring programs in addition to improvements in the modeling of the mass loss mechanism and the subsequent radiation (\\citet{Whitelock}, hereafter W94; \\citet{Trams  Obs} and \\citet{Jacco Mdot}, hereafter L99). These studies are essential because it is necessary to characterize the emission from individual AGB stars in order to understand the collective AGB emission from integrated galaxy light.  There are a number of key issues that IRAS was not able to address effectively due to instrumental limitations. First, it is important to definitively detect and characterize this MIR excess. Specifically, is there a spectral excess that peaks around 9.7\\( \\mu  \\)m and is consistent with a physical source of an AGB star plus circumstellar dust envelope? Second, it is useful to determine if the excess emission is similar from galaxy to galaxy, as would be expected for similar coeval populations. If there are significant differences, this emission could be used as a diagnostic tool for star formation or history, metallicity and age of these populations. Finally, it is important to determine if this excess follows a de Vaucouleurs law, as would be expected if the circumstellar envelopes around dying stars are the only source of the MIR excess \\citep{de Vaucouleurs}. We have used the CAM instrument on ISO to address these issues because it offers: 1) excellent sensitivity in six narrow-band filters that cover the 6-18\\( \\mu  \\)m region 2) high spectral resolution imaging using the CVF in the same spectral region and 3) spatial resolution of 6\\( \\arcsec  \\) over a field of view of 3.2\\( \\arcmin  \\). Here we report upon and analyze ISOCAM observations for nine nearby early-type galaxies.    In Section \\ref{Observations} we describe the observations made by ISO with the CAM narrow band filters and CVF.  We report on the extensive measures we undertook to process and understand the CAM detector data in Section \\ref{data}.  The mid-IR emission detected in these nine galaxies is described and characterized in Section \\ref{results}.  In Sections \\ref{E12} through \\ref{comparemdot} the mid-IR excess emission is compared to that of Galactic and LMC AGB stars to derive a scaling relation, revealing cumulative mass loss rates for the observed galaxies.  This observed mass loss rate is reasonably matched to theoretical predictions as described in Section \\ref{theory}.  Section \\ref{Light Profile Section} examines the radiation vs radius profile of these galaxies in different wavelength regions. We summarize and add final thoughts in Section \\ref{conclusion}. ",
        "conclusions": "Summary and Concluding Remarks}  Nine elliptical galaxies were observed with CAM on ISO in six narrow bands between 6\\( \\mu  \\)m and 15\\( \\mu  \\)m. From 6\\( \\mu  \\)m to 9\\( \\mu  \\)m the emission is consistent with the combined stellar emission from the K and M-type stars that dominate elliptical galaxy's integrated light. In eight of these galaxies we detected excess over stellar photospheric emission from 9\\( \\mu  \\)m to 15\\( \\mu  \\)m. For one galaxy, NGC 1404, ISOCAM CVF data, with its finer spectral resolution, shows the excess emission is consistent with the known 9.7\\( \\mu  \\)m oxygen-rich AGB silicate feature. We used Galactic and LMC AGB stars to calibrate a scaling relation, revealing galactic-wide mass loss rates for these galaxies. These observed rates mostly agree with theoretical predictions.  The observed rates do not scale with luminosity to a universal rate and thus suggests physical differences between these populations. We also show that emission at all wavelengths is consistent with a de Vaucouleurs' \\( R^{1/4} \\) law.   Now that it is possible to observe the signatures of mass loss, it would be valuable to revisit the predicted mass loss characteristics of cluster and galaxy sized populations. Tracking the population's total mass loss through time in a quantitative manner would be a valuable contribution to the field. Also the calibration of the relationship between mass loss and mid IR excess needs further investigation with additional studies of individual AGB stars. In particular, the effects of different metallicites should be investigated, as it is known to play a role, but the exact nature is not yet determined. An exciting avenue of research that will surely develop once the calibrations of these AGB features is accurately known, is the Visual-MIR color-color diagnostics.  This observational tool was recently explored in a theoretical study by \\citet{Bressan98}, which shows that for certain populations these types of relations could potentially break the age-metallicity degeneracy that has long plagued optical color-color relations. With SOFIA, SIRTIF and other upcoming IR missions, in addition to the continuing work though high altitude ground-based windows, it should be possible to continue to investigate these issues."
    },
    "0201/astro-ph0201430_arXiv.txt": {
        "abstract": "We describe a simple, efficient, robust and fully automatic algorithm for the determination of a Multi-Gaussian Expansion (MGE) fit to galaxy images, to be used as a parametrization for the galaxy stellar surface brightness. In most cases the least-squares solution found by this method essentially corresponds to the \\emph{minimax}, constant relative error, MGE approximation of the galaxy surface brightness, with the chosen number of Gaussians. The algorithm is well suited to be used with multiple resolution images (e.g., \\emph{Hubble Space Telescope} [HST] and ground-based). It works orders of magnitude faster and is more accurate than currently available methods. An alternative, more computing intensive, fully linear algorithm, that is \\emph{guaranteed} to converge to the smallest $\\chi^2$ solution, is also discussed. Examples of MGE fits are presented for objects with HST or ground-based photometry, including galaxies with significant isophote twist. ",
        "introduction": "\\footnotetext[3]{Based on observations made with the NASA/ESA {\\it Hubble Space Telescope}, obtained from the Data Archive at the Space Telescope Science Institute, which is operated by the AURA, Inc., under NASA contract NAS 5-26555.}  \\setcounter{footnote}{4}  Digital images of galaxies which are small by today's standards, may still contain over one million pixels. It is natural to try to distill the information contained in this large number of pixels into a set of quantities that can be used easily and efficiently  to derive information on the physical properties of the observed object. The most common approach currently consists in fitting ellipses of increasing semimajor axis to the galaxy images (e.g., Carter 1978; Kent 1984; Lauer 1985; Jedrzejewski 1987; Franx, Illingworth \\& Heckman 1989; Peletier et al.\\ 1990). Intensity profiles, ellipses shape and deviation from ellipses, parametrized by a few Fourier terms, are obtained with these methods. A nice feature of this parametrization is that it gives a physically meaningful description of the galaxy photometry in terms of the ellipticity, diskyness and boxyness at each radius.  One limitation of the ellipse fitting methods however is that strong deviations of the isophotes from ellipses cannot be easily modeled with a small number of Fourier terms. This makes these methods not well suited to the photometric modeling of multicomponent objects such as lenticulars and spirals, which have a bulge, a main disk and often an embedded nuclear disk or a bar. More importantly, if the deviations from ellipses are non negligible it becomes non trivial to use the fitted isophotes for the construction of realistic galaxy dynamical models due to the complexity of deprojection and the evaluation of the gravitational potential.  In this paper we focus on a completely different approach to the photometric modeling, the Multi-Gaussian Expansion (MGE) method, that overcomes both of these limitations. The original idea of the application of the MGE method to the problem of the deconvolution and deprojection of galaxy images comes from Bendinelli (1991). This idea was generalized to the non-spherical case and made applicable to real galaxies by Monnet, Bacon \\& Emsellem (1992) and further developed by Emsellem, Monnet \\& Bacon (1994a) and Emsellem, Dejonge \\& Bacon (1999).  The MGE method consists of a series expansion of galaxy images using two-dimensional (2D) Gaussian functions. The Gaussians have the beneficial properties that both the convolution (e.g., to take seeing or PSF effects into account) and the deprojection (to derive  the intrinsic stellar luminosity density from the observed galaxy photometry) can be performed analytically in a simple and efficient manner. As shown by Emsellem et al.\\ (1994a), many other dynamical and photometric quantities can be evaluated easily and accurately when the density is expressed in MGE form. For example, the MGE potential can be computed with a single integration, as opposed to the two that are required when the intrinsic density is stratified on similar triaxial ellipsoids, and three in the general case. In the case of $f(E,L_z)$ axisymmetric MGE dynamical models the velocity moments predicted from the Jeans equations, already projected onto the sky, can be expressed with only a double integration. The even part of the distribution function can also be easily retrieved from an MGE density distribution via the Hunter \\& Qian (1993) formalism.  The MGE parametrization is one of the few simple parame\\-trizations that are general enough to reproduce the surface brightness of realistic multicomponent objects (e.g., spirals with multiple disks). It has already been used for the modeling of a number of galaxies (Emsellem et al.\\ 1994b, 1996, 1999; Emsellem 1995; van den Bosch et al.\\ 1998; van den Bosch \\& Emsellem 1998; Cretton \\& van den Bosch 1999).  What is however still missing from the MGE machinery is an accurate, easy to use, generally available, robust and automatic algorithm for the determination of an MGE expansion from galaxy images. In this paper we present such an algorithm. The method we describe here brings the MGE fitting phase of a dynamical modeling process to the same level of the LOSVD extraction phase (e.g., van der Marel \\& Franx 1993) in terms of robustness and ease of use.  This paper is organized as follows. In Section~2 we define our notation and give some useful formulae for the application of MGE models to the dynamical modeling of galaxies. In Section~3 we discuss our new MGE fitting method. In Section~4 we give information on the availability of the software implementing the methods discussed in this paper. Some conclusions are given in Section~5. ",
        "conclusions": "The MGE method is one of the few simple parametrizations that are general enough to reproduce the surface brightness of realistic multicomponent galaxies. In addition many dynamical and photometric quantities can be evaluated easily and accurately when the density is expressed in MGE form.  In this paper we have described a simple yet powerful algorithm that reduces the process of generating an MGE fit to multiple galaxy images to a simple, fast and automatic task. We have provided examples of its practical use and have tested it by accurately reproducing the photometry of a relatively large sample of galaxies, both with and without isophote twists. We have also compared some of our fits with previously obtained photometric models.  We have also described an alternative algorithm that, although currently less practical, due to the larger computing power requirements, is guaranteed to converge to the minimum $\\chi^2$ solution within the accuracy imposed by an adopted grid of parameters in the solution space.  These algorithms have been implemented in an IDL program that can produce an MGE model starting from the observed images of a galaxy, requiring the user to only input the coordinates of the galaxy centre, the PA and a characteristic flattening. Multiple resolution images (e.g., ground-based and HST) can easily be fitted together, in one single step. The complete IDL source code implementing the algorithms described in this paper is made publicly available."
    },
    "0201/astro-ph0201405.txt": {
        "abstract": "The aim of these lecture notes is to familiarize graduate students and beginning postgraduates with the basic ideas of linear cosmological perturbation theory and of structure formation scenarios. We present both the Newtonian and the general relativistic approaches, derive the key equations and then apply them to a number of characteristic cases. The gauge problem in cosmology and ways to circumvent it are also discussed. We outline the basic framework of the baryonic and the non-baryonic structure formation scenarios and point out their strengths and shortcomings. Fundamental concepts, such as the Jeans length, Silk damping and collisionless dissipation, are highlighted and the underlying mathematics are presented in a simple and straightforward manner. ",
        "introduction": "%%%%%%%%%%%%%%%%%%%%%% Looking up into the night sky we see structure everywhere. Star clusters, galaxies, galaxy clusters, superclusters and voids are evidence that on small and moderate scales, that is up to 10~Mpc, our universe is very lumpy. As we move to larger and larger scales, however, the universe seems to smooth out. This is evidenced by the isotropy of the x-ray background, the number counts of radio sources and, of course, by the high isotropy of the Cosmic Microwave Background (CMB) radiation. The latter also provides a fossil record of our observable universe when it was roughly $10^5$ years old and about $10^3$ times smaller than today. So, the universe was very smooth at early times and it is very lumpy now. How did this happen? Although the details are still elusive, cosmologists believe that the reason is ``gravitational instability''. Small fluctuations in the density of the primeval cosmic fluid that grew gravitationally into the galaxies, the clusters and the voids we observe today. The idea of gravitational instability is not new. It was first introduced in the early 1900s by Jeans, who showed that a homogeneous and isotropic fluid is unstable to small perturbations in its density~\\cite{J}. What Jeans demonstrated was that density inhomogeneities grow in time when the pressure support is weak compared to the gravitational pull. In retrospect, this is not surprising given that gravity is always attractive. As long as pressure is negligible, an overdense region will keep accreting material from its surroundings, becoming increasingly unstable until it eventually collapses into a gravitationally bound object. Jeans, however, applied his analysis to a static Newtonian fluid in an attempt to understand the formation of planets and stars. In modern cosmology we need to account for the expansion of the universe as well as for general relativistic effects.  Despite the lack, as yet, of a detailed scenario, the rather simple idea that the observed structure in our universe has resulted from the gravitational amplification of weak primordial fluctuations seems to work remarkably well. These small perturbations grew slowly over time until they were strong enough to separate from the background expansion, turn around, and collapse into gravitationally bound systems like galaxies and galaxy clusters. As long as these inhomogeneities are small they can be studied by the linear perturbation theory. A great advantage of the linear regime is that the different perturbative modes evolve independently and therefore can be treated separately. In this respect, it is natural to divide the analysis of cosmological perturbations into two regimes. The early phase, when the perturbation is still outside the horizon, and the late time regime when the mode is inside the Hubble radius. In the first case microphysical processes, such as pressure effects for example, are negligible and the evolution of the perturbation is basically kinematic. After the mode has entered the horizon, however, one can no longer disregard microphysics and damping effects.  In these lectures we will a priori assume the existence of small inhomogeneities at some initial time in the early universe. A cosmological model is not complete, however, unless it can also produce these seed fluctuations through some viable physical process. Inflation (Guth (1981); Linde (1982)~\\cite{G}) appears to be our best option as it naturally produces a spectrum of scale-invariant gaussian perturbations. Topological defects (Kibble (1976)~\\cite{K}), such as cosmic strings, global monopoles and textures, offer a radically different paradigm to inflation for structure formation purposes. They have fallen out of favor, however, as their observational situation looks unpromising.  Understanding the details of structure formation requires, among other things, knowledge of the initial data. Structure formation, or galaxy formation as it is sometimes referred to, began effectively with the end of the radiation era at matter-radiation equality ($t_{\\rm eq}\\simeq4.4\\times10^{10}(\\Omega h^2)^{-2}\\,{\\rm sec}$, where $\\Omega$ is the density parameter of the universe). Thus, the start of the matter era also signals the beginning of structure formation. If we were ever to find out the details of how the structure in our universe formed we need to know the initial data at that epoch. The necessary information includes: (i) the total amount of the non-relativistic matter; (ii) the composition of the universe and the contribution of its various components to the total density, namely $\\Omega_{\\rm b}$ from baryons, $\\Omega_{\\gamma}$ from relativistic particles, $\\Omega_{\\rm WIMP}$ from relic WIMPs,\\footnote{Weakly Interacting Massive Particles (WIMPs) are stable non-baryonic species left over from the earliest moments of the universe.} $\\Omega_{\\Lambda}$ from a potential cosmological constant etc; (iii) the spectrum and the type (i.e.~adiabatic or isothermal) of the primeval density perturbations. Given these one can, in principle, construct a detailed scenario of structure formation, which then will be tested against observations. The importance of specifying the initial conditions is paramount, since inverting present observations to infer the initial data is unfeasible after all the astrophysical filtering that has taken place. Speculation on the history of the early universe, backed by recent observations provide some ``hints'' as to the appropriate initial data. They point towards $\\Omega=1$ from inflation; $\\Omega_{\\rm WIMP}\\simeq0.3$ (including a small baryonic contribution) and $\\Omega_{\\Lambda}\\simeq0.7$ from nucleosynthesis, inflation and the supernovae redshift measurements; and adiabatic fluctuations with a Harrison-Zeldovich spectrum~\\cite{H} from inflation.  The layout of these notes is as follows. In Sec.~2 we present the Newtonian treatment of linear cosmological perturbations, discuss issues such as pressure support and the ``Jeans length'' and provide the key results. The general relativistic analysis is outlined in Sec.~3 and the basic linear equations are derived. We also give a brief discussion of the ``gauge problem'' and provide characteristic solutions of the relativistic approach. In Sec.~4 we discuss entirely baryonic structure formation scenarios, emphasizing the collisional damping of adiabatic density perturbations. Non-baryonic ``hot'' and ``cold'' dark matter models are presented in Sec.~5, together with their advantages and shortcomings. The aim of these lectures notes is to provide the basic background to graduate students in physics and astronomy as well as to beginning postgraduates. We would like to familiarize the newcomer with fundamental concepts such as gravitational instability, the Jeans length, collisional and collisionless damping. The necessary mathematics are also provided in simple and straightforward manner. Overall we want to give a brief but comprehensive picture of the linear regime, so that the interested student will feel more confident when looking at more sophisticated treatments. For further details we refer the reader to some of excellent monographs that now exist in the literature (see~\\cite{KT} for a list of them). The lectures do not require a particularly specialized background, although some knowledge of cosmology and general relativity will be helpful.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "%%%%%%%%%%%%%%%%%%%% The question of how the observed large-scale structure of the universe developed and how galaxies were formed has been one of the outstanding problems in modern cosmology. Looking back into the past hundred years one sees three decisive moments in the pursuit of the answer. The first milestone was the formulation of general relativity which provided researchers with the theoretical tool to probe the large-scale properties of the universe. Hubble's observations manifesting the expansion of the universe was the second decisive moment, as it forced cosmologists to break away from the then prevailing concepts of a static and never-changing cosmos. Finally, the discovery of the Cosmic Microwave Background radiation by Penzias and Wilson established the Hot Big Bang theory, an idea advocated several years earlier primarily by Gamov. At last, cosmologists had a definite model within which they could tackle the structure formation question. Pending future observations, one could argue that the inflationary paradigm is one additional milestone in our effort to understand the workings of our universe. The recent supernovae measurements, which suggest an accelerated universal expansion, could prove another very decisive moment. Time will show whether they actually are.  Despite the problems and the uncertainties, cosmologists now believe that all the structure that we observe around us today originated from minute perturbations in a cosmic fluid that was smooth to the accuracy of one part in ten thousand at the time of recombination. Such tiny irregularities could have been triggered by quantum fluctuations that were stretched out during the inflationary expansion or by topological defects such as cosmic strings for example. Given the current observational status, inflation appears to be the strongest candidate. In these notes we have set aside the question of the origin of the primeval fluctuations. Our discussion focussed on the linear evolution of these minute irregularities once the universe entered the post-inflationary Hot Big Bang era, and on the physical processes that could have affected them. When studying density perturbations, one sooner or later encounters the reality that baryon inhomogeneities cannot actually grow before recombination. This fact, together with the extreme smoothness of the CMB temperature, implies that baryonic fluctuations simply do not have enough time to produce the observed structure. When one adds to that the nucleosythesis constraints and the strong theoretical and observational bias for spatial flatness, the once popular baryon-dominated picture of the universe seems unsustainable. So, if the baryons are not the dominant form of matter in our universe, then what is? The answer to that might lie in high energy physics theories. Theoretical physics provides a whole zoo of supersymmetric, dark matter species that could bring $\\Omega$ close to unity and also ``assist'' structure formation. The attractive feature of collisionless matter is that perturbations in its density start growing earlier than those in the baryonic component. Thus, as soon as the baryons decouple from radiation, they undergo a period of fast growth as they fall into the potential wells of the collisionless species. This can improve the final picture but unfortunately does not solve all the problems. The key obstacle being that a single collisionless species does not seem capable of fitting all the data. At best, one needs the presence of one cold and one hot dark matter species in a combination that will make the most of their advantages while minimizing their shortcomings. The recent supernovae results, suggesting that our universe might be dominated by some sort of dark energy have added extra flavor to the whole situation. The presence of a cosmological constant, or quintessence, means that researchers have extra freedom when dealing with crucial cosmological parameters, such as the age of the universe for example. On the other hand, however, the idea that the majority of the matter in our universe is in the form of some unknown exotic species brings back some rather embarrassing memories from our relatively recent past. A lot could be decided within the next few years as we expect an influx of high quality data. The ``Boomerang'' and ``Maxima'' observations have added valuable information which seems to favor the inflation based models. The near future ``Probe'' and ``Planck'' satellite missions also promise high precision data. For some researchers structure formation is a story that is fast reaching its conclusion. The future will show if they are right or just too hasty."
    },
    "0201/astro-ph0201095_arXiv.txt": {
        "abstract": "We use numerical  simulations of structure formation in  a Cold Dark Matter (CDM) cosmology to compare the angular momentum distributions of dark matter and non-radiative gas in a large sample of halos.  We show  that   the  two  components  have   identical  spin  parameter distributions and  that their angular  momentum distributions within individual halos  are very similar, all in  excellent agreement with standard assumptions.  Despite  these similarities, however, we find that the  angular momentum  vectors of the  gas and dark  matter are poorly  aligned,  with a  median  misalignment  angle  of $\\sim  30$ degrees,   which  might   have  important   implications   for  spin correlation  statistics used  in  weak lensing  studies. We  present distributions  for the  component of  the angular  momentum  that is aligned with the total angular  momentum of each halo, and find that for between 5 and 50 percent of the mass this component is negative. This  disagrees  with  the  generally  adopted  `Universal'  angular momentum  distribution, for  which the  mass fraction  with negative specific angular  momentum is zero.  We discuss  the implications of our  results for  the formation  of disk  galaxies.   Since galactic disks generally  do not contain counter-rotating stars  or gas, disk formation  cannot  occur  under  detailed conservation  of  specific angular  momentum.   We  suggest  that the  material  with  negative specific  angular momentum combines  with positive  angular momentum material  to  build a  bulge  component, and  show  that  in such  a scenario  the remaining  material can  form  a disk  with a  density distribution that is very close to exponential. ",
        "introduction": "\\label{sec:intro}  Understanding  the  structure  and   formation  of  disk  galaxies  is intimately linked to understanding the origin of its angular momentum. In hierarchical  structure formation  scenarios the luminous  parts of galaxies  form from  gas that  is cooling  and condensing  within dark matter (DM) halos  and these merge to build  larger and larger objects (White \\& Rees  1978).  Within this picture, the  current paradigm for disk formation  contains three important ingredients:  (i) the angular momentum originates  from cosmological torques (Hoyle  1953), (ii) the gas  and dark matter  within virialized  systems have  initial angular momentum distributions  that are identical (Fall  \\& Efstathiou 1980), and (iii) the gas conserves its specific angular momentum when cooling (Mestel 1963).  It  is  well  established  that cosmological  torques  impart  angular momentum  to dark matter  halos.  Many  studies have  investigated the resulting  angular momenta  of dark  matter halos  using  either tidal torque  theory (e.g.,  Peebles  1969; Doroshkevich  1970; White  1984; Catelan  \\&  Theuns 1996)  or  N--body  simulations  (e.g., Barnes  \\& Efstathiou 1987; Warren \\etal  1992; Bullock \\etal 2001). The specific angular  momentum   of  halos   is  typically  parameterized   by  the dimensionless spin parameter \\begin{equation} \\label{lam} \\lambda \\equiv \\frac{|{\\bf J}| \\, |E|^{1/2}}{G \\, M^{5/2}}, \\end{equation} with $G$  the gravitational constant and  ${\\bf J}$, $E$,  and $M$ the total angular  momentum, energy and  mass of the  object, respectively (Peebles 1969).  Under the assumptions (ii) and (iii) disks have scale lengths $R_d \\propto \\lambda R_{\\rm vir}$.  Given the distributions of halo  spin parameters  and of  halo  virial radii  $R_{\\rm vir}$,  the implied distribution  of disk scale lengths is  in excellent agreement with observations  (e.g., Dalcanton, Spergel \\& Summers  1997; Mo, Mao \\& White 1998; de Jong \\& Lacey 2000).  This success has prompted many detailed  studies of  disk galaxy  formation, always  under  the three assumptions  listed  above (van  den  Bosch  1998,  2000, 2001,  2002; Jimenez \\etal  1998; Natarajan 1999; Heavens \\&  Jimenez 1999; Firmani \\& Avila-Reese  2000; Avila-Reese \\& Firmani  2000; Buchalter, Jimenez \\& Kamionkowski 2001).  The standard  picture of  disk formation that  has emerged  from these studies has been remarkably successful in explaining a wide variety of observational properties  of disk galaxies.   However, two significant problems, both  related to the  angular momentum, have come  to light. First of  all, detailed hydro-dynamical simulations  of disk formation in a cold dark matter (CDM)  Universe yield disks that are an order of magnitude  too small  (Navarro \\&  Benz 1991;  Navarro \\&  White 1994; Steinmetz \\& Navarro 1999;  Navarro \\& Steinmetz 2000).  This problem, known as  the angular  momentum catastrophe, is  a consequence  of the hierarchical formation of galaxies which  causes the baryons to lose a large fraction  of their angular momentum  to the dark  matter.  It is unclear whether  this is a  problem for the  theory or merely  for the particular  way  in  which  feedback  processes  are  implemented  (or ignored) in the simulations (e.g., Navarro \\& White 1994; Weil, Eke \\& Efstathiou  1998;  Sommer-Larsen, Gelato  \\&  Vedel  1999; Thacker  \\& Couchman 2001).  The  second problem concerns  the actual  density structure  of galaxy disks.    Assuming  detailed   angular  momentum   conservation,  this structure  is a  direct  reflection of  the  distribution of  specific angular momentum in the proto-galaxy.  Bullock \\etal (2001) determined the angular  momentum distributions  of individual dark  matter halos, which  according to assumption  (ii) should  also reflect  the angular momentum distribution of the gas  that forms the disk.  However, these distributions seem to have far  too much low angular momentum material to be consistent with the typical exponential density distributions of disk galaxies (Bullock \\etal 2001;  van den Bosch 2001; van den Bosch, Burkert \\& Swaters 2001). Knebe  \\etal (2002) and Bullock, Kravtsov \\& Col\\'{\\i}n (2002)  investigated the angular momentum  profiles of warm dark  matter (WDM)  halos.   Except  for a  slight  difference in  the average spin  parameter, CDM and  WDM halos have very  similar angular momentum distributions,  implying that this  particular problem cannot be solved  by invoking WDM instead  of CDM. A  more promising solution has  recently been  suggested  by  Maller \\&  Dekel  (2002). In  their picture  halo angular  momentum  originates from  the orbital  angular momentum of  merging satellites, rather  than from pure  tidal torques (see also  Maller, Dekel \\&  Somerville 2001). Combining  this picture with a simple  model for feedback and tidal  stripping Maller \\& Dekel are able  to explain the  angular momentum profiles of  dwarf galaxies measured by van den Bosch \\etal (2001).  These angular  momentum problems seem  to imply that  assumptions (ii) and (iii) listed  above, and which have been  essential ingredients of our  standard paradigm  for the  formation  of disk  galaxies, may  be incorrect.  Remarkably, despite many papers that have investigated the angular momentum  distributions within dark matter halos,  to the best of our knowledge, no study  so far has addressed the spin distribution of the  gas in these halos in  order to test whether  the {\\it ansatz} that the  gas and  dark matter have  similar initial  distributions of specific  angular momentum is  correct.  In  fact, it  seems plausible that they  may not.  Gas  and dark matter suffer  different relaxation mechanisms  during halo  collapse: Whereas  the dark  matter undergoes collisionless  virialization  through   violent  relaxation,  the  gas attempts  to achieve  a  hydrostatic equilibrium  through shocks.   In order to  test whether the  initial angular momentum  distributions of gas  and  dark  matter   are  indeed  similar,  we  perform  numerical simulations  of  structure   formation  in  a  $\\Lambda$CDM  cosmology including both dark matter  and a non-radiative gas.  After presenting our    simulations     and    outlining    our     analysis    methods (Section~\\ref{sec:method}),  we  present  various  statistics  of  the angular   momentum   distributions  (Section~\\ref{sec:results}).    In Section~\\ref{sec:diskform} we discuss  the implications of our results for the theory of disk  galaxy formation, and we summarize our results in Section~\\ref{sec:concl}. ",
        "conclusions": "\\label{sec:concl}  It  is  generally assumed  that  the  gas and  the  dark  matter in  a proto-galaxy have the same  distribution of specific angular momentum. This is motivated by the fact that both mass components experience the same  tidal  torques.   However,  the two  components  undergo  rather different  relaxation mechanisms:  whereas the  dark  matter undergoes collisionless, violent relaxation, the gas gets shocked which could in principle cause a redistribution of its angular momentum distribution. In  this paper we  have used  numerical simulations  to test,  for the first time,  whether the  {\\it ansatz} that  gas and dark  matter have similar  initial  angular   momentum  distributions  is  correct.   We presented  a   numerical  simulation  of  structure   formation  in  a $\\Lambda$CDM cosmology including both  dark matter and a non-radiative gas.   For $~\\sim  380$ halos  with virial  masses in  the  range $4.8 \\times  10^9 h^{-1}  \\Msun \\leq  M_{\\rm vir}  \\leq 1.1  \\times 10^{12} h^{-1} \\Msun$ (at  $z=3$) we computed the angular  momenta of both the gas and the dark matter.  We have  shown that on average the  gas and dark matter  have the same distribution  of  spin  parameters  and that  their  detailed  angular momentum  distributions in  individual halos  are  remarkably similar. Evidently, shocking during the virialization process does not decouple the  angular  momentum  of the  gas  from  that  of the  dark  matter, supporting  the standard  assumption  that gas  and  dark matter  have identical  angular momentum  distributions. We  have shown  that these results  are robust  by  also analyzing  a  high resolution  numerical simulation of  a single cluster-sized halo  at $z=0$ in  a much larger cosmological volume.  A  careful analysis  of  the detailed  angular momentum  distributions reveals  that  they  are  inconsistent with  the  `Universal'  profile suggested  by B01.   In the  angular momentum  distributions presented here  between 5  and  50 percent  of  the mass  has negative  specific angular momentum, whereas B01  have (implicitly) assumed that all mass has  positive  specific angular  momentum.   Furthermore, our  angular momentum distributions  have wings out  to very high  specific angular momentum,  whereas the profile  suggested by  B01 has  a cut-off  at a certain  $j_{\\rm max}$  that results  in disk  truncation at  too high surface density (van den Bosch 2001).  The fact  that relatively large mass fractions  have negative specific angular  momentum suggests  that the  standard assumption  of detailed specific  angular  momentum  conservation,  which  is  often  made  in theories  of  disk formation,  cannot  be  correct.   After all,  disk galaxies  do not  typically  contain material  with negative  specific angular momentum.   We have suggested that during  the cooling process the gas with negative specific angular momentum may collide/shock with material  with positive  specific angular  momentum to  build  a bulge component with zero  angular momentum. In such a  picture at least the {\\it total} amount of angular momentum of the gas is conserved, and we have  shown that it  can produce  disks with  near-exponential surface density distributions  without the  small truncation radii  implied by the  `Universal'  profile  of  B01.   Numerical  simulations  of  disk formation that include cooling do  indeed seem to reveal the formation of a dense central knot of gas  in addition to the disk (e.g., Katz \\& Gunn 1991; Navarro \\& White 1994).  Although this `bulge' is typically interpreted  as forming  out of  sub-clumps that  lose  their specific angular momentum  to the dark  matter through dynamical  friction, our results  suggest  that even  without  substructure  a  bulge may  form naturally.  In  fact, taking the  results at face value  suggests that $\\sim  40$ percent  of all  halos are  unable to  form  disk dominated galaxies,  simply because  they  have too  large  mass fractions  with negative  specific  angular   momentum.   In  addition,  systems  with bulge-to-disk mass ratios below 10 percent will be extremely rare.  In reality the formation of  disk galaxies will be more complicated as depicted above.   We have shown  that the angular momentum  vectors of the gas  and dark  matter are misaligned  by on average  $\\sim 35^{\\rm o}$. Surprisingly the amount of  misalignment is not a measure for the degree to which the respective angular momentum distributions agree or disagree  with  each  other.   There   is  a  weak  tendency  for  the misalignment  to be  weaker in  halos that  are more  massive  or have larger spin parameters. Higher  resolution simulations are required to investigate the amount and origin of this misalignment in more detail. In  addition to  this  misalignment between  the  two components,  the angular  momentum  vector  of  each individual  component  can  change direction   quite  dramatically  with   radius.   This   is  typically interpreted as  originating from  decoherence in the  angular momentum orientation of material accreted at different times (Ryden 1988; Quinn \\& Binney 1992) and is often invoked as a mechanism to create warps in disk  galaxies (e.g., Ostriker  \\& Binney  1989).  Although  the exact implications of these various misalignments are currently unclear, and require  a   more  detailed  investigation,  it  is   clear  that  any misalignment of the gas with  either itself or with the inertia moment of the dark matter  halo has to vanish if one wants  to build a stable disk.  If this  results in a significant transfer  of angular momentum from  the gas  to  the dark  matter  these misalignments  might be  an additional cause  for the  angular momentum catastrophe  hampering the formation  of   sufficiently  extended  disk   galaxies  in  numerical simulations.  Finally, we  emphasize that the implications  for disk/bulge formation discussed above  are highly speculative.   In our simulations  the gas was  not   allowed  to  cool,   and  we  used  the   angular  momentum distributions of the gas found at $z=3$ to speculate about the fate of the gas once cooling were turned  on.  In reality the gas will already have  (partially) cooled  inside  sub-clumps that  merge  to form  the systems  analyzed here,  and  this  will also  impact  on the  angular momentum distribution  of the gas and its  subsequent evolution (e.g., dynamical friction).  On the other hand, it has been argued by several authors that in order to  prevent the angular momentum catastrophe the gas has to be prevented from cooling inside sub-clumps for instance by feedback processes  or preheating (e.g., Navarro \\&  White 1994; Weil, Eke \\&  Efstathiou 1998; Sommer-Larsen, Gelato \\&  Vedel 1999).  Under those circumstances  the angular momentum distribution of  the gas may in fact resemble  those analyzed here, although it  is unclear to what extent  these   heating  processes  influence   the  angular  momentum distributions of  the gas  prior to being  incorporated in  the halos. For instance,  processes related to feedback  and re-ionization create pressure gradients in the gas  which may decouple the angular momentum distribution of the  gas from that of the dark  matter. In a follow-up paper  (Abel \\etal 2002)  we address  in more  detail the  impact such pressure  forces may  have on  the angular  momentum  distributions in proto-galaxies."
    },
    "0201/astro-ph0201210_arXiv.txt": {
        "abstract": "The Cannonball Model is based on the hypothesis that GRBs and their afterglows are made in supernova explosions by relativistic ejecta similar to the ones observed in quasars and microquasars. Its predictions are simple, and analytical in fair approximations. The model describes well the properties of the $\\gamma$-rays of GRBs. It gives a very simple and extremely successful description of the optical and X-ray afterglows of {\\it all} GRBs of known redshift. The only problem the model has, so far, is that it is contrary to staunch orthodox beliefs. ",
        "introduction": "The idea that GRBs are due to collimated emissions is not recent. In the case of GRBs from quasars it was discussed by Brainerd \\cite{Bra}; in the case of a funnel in an explosion, by Meszaros \\& Rees \\cite{MR}. In what is no doubt the relevant case: jets in stellar gravitational collapses, the idea has been developed over the years by Dar and collaborators \\cite{SD} to \\cite{DD2001c}. Now we know that long-duration GRBs are cosmological, originate in galaxies, are associated with supernovae (SNe) and have energies that would be ridiculously large for a stellar spherical explosion (a fireball). GRBs must be ``jetted''.    In the currently dominating scenarios, the ejecta that beget a GRB and its afterglow (AG) are thrown off in a uniform cone with an opening angle $\\theta_j(0)$. This cone expands sideways: the angle $\\theta_j(t)$ subtended by the ejecta {\\it increases} with time, delineating a {\\it firetrumpet}, as in Fig.~(1). Relativistic jets are ubiquitous in astrophysics. The ejecta of these real jets, as seen from their emission point up to the point where they eventually stop and expand, generally subtend angles that {\\it decrease} with time: just the opposite of the assumed behaviour of firetrumpets. In the analysis of the observed jets, e.g. \\cite{PZ,RM,GC}, it is the fixed angle of observation ---and not the angle subtended by the ejecta--- that plays a key role.  \\begin{figure} \\includegraphics[height=.5\\textheight]{coneCB} \\caption{(a) A firecone or, more properly, a {\\it firetrumpet}. In these scenarios the cone expands conically for a distance, after which the jet angle $\\rm\\theta_j$ widens faster as its firefront travels. (b) Cannonballs (shown here, somewhat pedantically, a bit Lorentz-contracted) subtend decreasing angles as they travel. The only relevant angle in the CB model is the observer's viewing angle $\\theta$.} \\end{figure}   The Cannonball (CB) model is based on the contention that GRBs and their AGs are made by relativistically jetted balls of ordinary ``baryonic'' matter which, by a mechanism \\cite{DD2001c} that I will outline, stop expanding soon after their emission. The CB idea gives a good description of the properties of the $\\gamma$-rays in a GRB, that we modelled in simple approximations \\cite{DD2000b}. It gives an excellent and complete description of optical and X-ray afterglows, which we have modelled in full detail, as I outline here. ",
        "conclusions": "At the time I gave this talk and submitted it to the Woods Hole GRB conference proceedings, we had not yet worked out the CB-model's predictions for radio afterglows. By the time I am posting it on the web, we have made significant progress: once again, the CB model's predictions are uncharacteristically simple and succesful \\cite{DDD2002}. We have no CB-model explanation for the scintillation behaviour of GRB 970508 \\cite{Ta}. Other than that, the CB model explains well all properties of GRBs. In the case of optical and X-ray afterglows, which I have outlined, the CB-model's predictions are univocal (as opposed to multiple-choice), very explicit, analytical in fair approximations, quite simple, very complete,  and extremely successful. I doubt that this statement applies to other models of GRBs.   I am not saying that the CB model will stand all future tests. But I would expect that ---confronted with a simple and successful model--- most scientists would, at least, say: {\\it Hum! } and ask good questions. Not the case. Four of our papers on this subject, \\cite{DD2000a} to \\cite{DD2001b}, have already been rejected by referees who found no single error, and/or stuck to bad numerical questions, even after being proved numerically (i.e. inarguably) wrong. This rejection statistics makes me feel that the GRB community has officially certified us... as crackpots. All this reflects, I suspect, the global rise of fundamentalism. Some people, almost literally in this case, still refuse to ``look through the telescope''. A GRB theorist, sitting on the first row in my talk, was kind enough to enact my social comments. Indeed, he rose in ire to exclaim: {\\it I am glad that the referee system is working!} He then stuck to a bad numerical question. Most of the rest of the questions period was not this aggressive\\footnote{Most... but not all. Another GRB theorist stated his suspicion that I was not presenting radio results, not because we had not yet worked them out, but because they were no good!!! For my first experience at a GRB conference, this was not bad: a bit like dancing flamenco on a mine field.}: there is still hope in science's sempiternal contest with faith.  \\clearpage \\newpage"
    },
    "0201/astro-ph0201026_arXiv.txt": {
        "abstract": "The convective collapse of thin magnetic flux tubes in the photospheres of sun-like stars is investigated using realistic models of the superadiabatic upper convection zone layers of these stars.  The strengths of convectively stable flux tubes are computed as a function of surface gravity and effective temperature. We find that while stars with T$_{eff}\\ge$ 5500 K and log $g$$\\ge$ 4.0 show flux tubes highly evacuated of gas, and hence strong field strengths, due to convective collapse, cooler stars exhibit flux tubes with lower field strengths.  Observations reveal the existence of field strengths close to thermal equipartition limits even in cooler stars, implying highly evacuated tubes, for which we suggest possible reasons. ",
        "introduction": "\\label{sec:intro}  An important physical quantity basic to the understanding of stellar magnetism is the field strength in small-scale flux tubes and spots on the photospheres. Measurements and interpretation of magnetic field strengths \\citep{saar88,saar96,guenther97} in cool main-sequence stars other than the Sun have been a subject of much debate recently \\citep{basrietal90, safier99}, because of the difficulties associated with the modeling of the atmospheric structural changes that the inhomogeneous fibril state of the magnetic field introduces \\citep{basrietal90}. However, recent improvements in observing techniques \\citep{saar96,donatietal97,j.krulletal99} have increased the reliability of observational determination of stellar magnetic field strengths. The observed Zeeman broadening on cool stars is believed to be produced by small flux tubes that appear bright, similar to the solar magnetic network and facular bright points, rather than spots, which are much fainter than quiet photospheres and hence contribute little to the stellar profiles \\citep{basrietal90,safier99}.  Based on measurements of Zeeman effect in cool main sequence stars \\citet{saar94,saar96} has inferred that the magnetic field strengths are close to the thermal equipartition field strengths B$_{eq}$=$\\sqrt{ 8\\pi p_{e}}$ at the observed levels in the atmosphere, where $p_{e}$ is the gas pressure in the unmagnetised atmosphere, with a conclusion that the surface distribution is in the form of highly evacuated small flux tubes in pressure balance with the ambient atmosphere. Various other measurements \\citep{saarlinsky85,basri-marcy94,j.krulletal99} also establish such field strengths leading to the general acceptance \\citep{linsky99} that stellar surface field strengths scale as the square root of the surface gas pressure. We note that B$_{eq}$, at any geometrical level in a stellar atmosphere, is the maximum possible field strength for flux tubes confined by gas pressure and can be easily determined \\citep{bunte-saar93,safier99} as a function of log $g$ and T$_{eff}$ using atmospheric models (e.g., \\citet{kurucz93}), where $g$ and T$_{eff}$ denote the surface gravity and effective temperature.  The widely accepted mechanism to account for the kG range field strengths in solar magnetic flux tubes is the superadiabatic effect or convective collapse (CC, hereafter) \\citep{parkr78,sptz79}.  In this Letter, we report results from a study of this mechanism in the stellar context and test its efficiency as a function of log $g$ and T$_{eff}$.  Concentration of magnetic fields into discrete flux tubes or sheaths has long been recognised as a general consequence of the interaction between magnetic fields and cellular convection that expels the flux into downdrafts \\citep{parkr63,galloway-weiss81,proctor-weiss82}.  High-resolution observations of solar surface magnetic fields confirm the operation of such flux expulsion process. The CC, which is a consequence of thermal insulation of the expelled magnetic flux against convective motions and the superadiabatic thermal stratification of the ambient atmosphere, further intensifies the field to observed values \\citep{parkr78}. It is a convective instability, modified by the magnetic field, and drives a down flow along the field lines of a flux concentration leading to the evacuation of gas inside; lateral pressure balance ensures that a highly compressed intense field flux tube or sheath is formed. This instability develops very rapidly, typically on a free-fall time scale, and is faster than other MHD instabilities present.  Hence, the criterion for convective stability can be regarded as a test to check for the occurrence of flux tubes in stars.  On the Sun, this instability has a typical growth time of about 2 - 5 minutes \\citep{hasan84,raj00}. Other important instabilities are the interchange instability \\citep{parkr75,piddn75} connected with the field topology, which has been studied in the stellar context by \\citet{bunte-saar93}, and the Kelvin- Helmholtz instability \\citep{schuss79,tsinga80} connected with the relative motion between the field confined and the surrounding gas.  We use model atmospheres constructed using the ATLAS9 model atmosphere code \\citep{kurucz93}, which can extend the \\citet{kurucz93} atmospheres to cover deeper regions of the convection zone, to study the CC of thin flux tubes embedded in them. The critical field strengths for the tubes to be stable, which indicate the amount of evacuation that stellar flux tubes undergo as a result of the collapse process, are found. ",
        "conclusions": "We have examined the superadiabatic effect (or CC) \\citep{parkr78} in the stellar context, using realistic models of the outer convective layers of stars. Our results show that whereas it is possible to produce highly evacuated stable tubes in stars with T$_{eff}\\ge$ 5500 K through the CC mechanism, it is not so in cooler stars; the decreasing amount of superadiabaticity in the upper convection zone layers of K and M spectral type stars make the CC inefficient yielding field strengths much less than B$_{eq}$. Hence, if CC is the main physical process responsible for the formation of fibril-like evacuated flux tubes then it is expected that the differences between the observed B and B$_{eq}$ increase as T$_{eff}$ decreases. The existing observational results are not entirely consistent with such a trend as there are several cases of K and M dwarfs having B values close to or even exceeding B$_{eq}$; but, all these stars with high B have high filling factors, $f$, typically 0.5 or higher, which correlate strongly with stellar angular rotation frequency $\\Omega$. Rotation-activity correlation dominated through high $f$ values is a well observed fact \\citep{saar90, saar96} consistent with dynamo theory predictions. A careful look at the observational results compiled by \\citet{saar90} and \\citet{solanki92} reveal that all the cases of B $<$ B$_{eq}$ also belong to K and M spectral types, but which have smaller $\\Omega$ and are less active; and all the G type stars show B values tightly around B$_{eq}$. As an example, from the compilation by \\citet{saar90}, we find that the stars GL 171.2A (BY Dra) and HD 201091, which are of similar spectral type (K5V, T$_{eff}$$\\approx$4400 K), have widely differing field strengths: 2.8 kG \\citep{saaretal87} and 1.2 kG \\citep{marcy-basri89} respectively. But these field strengths have good correlation with the $\\Omega$ and $f$ values (rotation periods of 1.85 and 37.9 days and $f$ values of 0.5 and 0.24, respectively). From our results shown in Figure 1, we find for T$_{eff}$$\\approx$4400 K and a log $g$$\\approx$4.5 (main-sequence) a field strength of $\\approx$ 1.3 kG that CC yields, in close agreement with the slow rotating, low $f$ value case of HD 201091.  \\citet{marcy-basri89} themselves caution that the separation of B and $f$ values is uncertain and that the flux $f$B is larger than expected for HD 201091. But, their speculation that B remains at the inferred value while $f$ changes by 2 orders of magnitude, which is not inconsistent with the observed time variation in the chromospheric emission, finds support from the present result.  The best support for the present calculations comes from a multi-line infra-red Zeeman analysis of $\\epsilon$ Eridani, a K2 V star with T$_{eff}$=5130 K and log $g$=4.7, by \\citet{valentietal95}, who find $f$=0.088 and B=1.44 kG in close agreement with the value B$_{c}\\approx $ 1.4 kG from Figure 1.  It would thus seem that the CC indeed operates on all solar-like main sequence stars producing convectively stable tubes as dictated by the sub-surface superadiabatic structure; and, the higher than expected B values in K and M dwarfs may originate from effects induced by high values of $\\Omega$.  Recent detections of increased photometric variability in highly active K dwarfs, which exhibit $saturation$ in their magnetic activity \\citep{O'delletal95}, indicate that stars with angular velocities $\\Omega >\\Omega_{sat}$ show increased number of spots. Thus, the present result that B$_{c}$ $\\ll$B$_{eq}$ for non-spot small scale magnetic fields in stars with T$_{eff}\\le$5000 K and its agreement with B observed on slow rotators provide a theoretical reason to believe the idea that $\\Omega_{sat}$ marks a change in contributions from $f$ and B to $f$B \\citep{saar96b}: $f$ saturates around 0.6 and $\\Omega$($ > \\Omega_{sat}$) begins to contribute to B by increasing $f_{spot}/f$ to maintain the continued increase of flux $f$B well described by the power-law fits $f$B$\\propto \\Omega^{1.3}$ or $f$B$\\propto (\\tau_{c}\\Omega)^{1.2}$ to the observations \\citep{saar90}, where $\\tau_{c}$ is the convective turn-over time and $\\tau_{c}\\Omega$ is the inverse Rossby number.  We suggest that the calculated stability limits B$_{c}$ can be used to separate $f$ from the observed fluxes $f$B for slow rotators ($\\Omega < \\Omega_{sat}$), when there remain otherwise unknown uncertainties in the separation of $f$ and B in observational analyses. The calculations here extend down only to T$_{eff}$=4000 K due to limitations in generating reliable convection zone models for cooler stars, but, by extrapolation, it is likely that CC remains ineffective yielding B$_{c}(\\tau_{c}=1)\\ll $ B$_{eq}(\\tau_{c}=1)$ in cooler M dwarfs too.  Two similar spectral type M dwarfs, observed by \\citet{krull-valenti96}, do show widely differing field strengths, in correlation with their $\\Omega$, but the lower $\\Omega$ star seems to show B larger than the likely B$_{c}$. Future refinements, both in observations and theory, are needed to understand fields in M dwarfs.  We note that the thin tube approximation does not get constrained by increasing $f$;  as long as the distribution of flux remains in the form of small flux tubes dominating the observed spectral profiles, the present results on B$_{c}$ can be used in the interpretation of observations. Thus, if the observed profiles do arise only from non-spot fields even in fast rotating K and M dwarfs which show B$\\gg$B$_{c}$ then the present results call for new physical mechanisms to produce highly evacuated tubes, which are likely not due to thermodynamic reasons but due to fast-rotation induced phenomena. This may imply a failure of solar analogy as far as the formation and dynamics of surface magnetic fields are concerned. We conclude by pointing out that, in any case, the convective stability limits calculated here need to be satisfied by pressure confined flux tubes and thus may serve as useful lower limits."
    },
    "0201/astro-ph0201356_arXiv.txt": {
        "abstract": "We present 236 new radial velocities of galaxies in the cluster A2256 measured with the WIYN Hydra multi-object spectrograph.  Combined with the previous work of \\citet{FAB89}, we have velocities for a total of 319 galaxies of which 277 are cluster members.  In addition to the new radial velocities, we present a $3 \\times 3$ image mosaic in the R band of the central $19\\arcmin \\times 19\\arcmin$ region of A2256, from which we obtained photometry for 861 galaxies.  These data provide strong evidence for a merger event between two groups.  In addition, we present evidence for the presence of a third group, on the outer reaches of the system, that is just now beginning to merge with the system. ",
        "introduction": "\\label{sec:introduction_A2256}  Originally thought to be a highly relaxed or evolved cluster, A2256 was shown by \\citet{DRESSLER78} to show evidence for equipartition in energy.  With a richness class 2, A2256 is surpassed only by 5\\% of the clusters found in Abell's (1958) catalog.  \\citet{FABER77} measured the velocity dispersion to be $1351\\ \\kmsec$, making it one of the highest velocity dispersions of known galaxy clusters.  \\citet{CARTER80} showed the galaxy distribution to be elliptical.  \\citet{FAB89} increased the number of measured radial velocities to 89 and found a flat velocity histogram.  They first suggested the presence of substructure, and interpreted this as indicating that A2256 represents a merger in progress.  In the X-ray, A2256 is a luminous and hot cluster with a central temperature of $\\sim\\!\\!7$\\kev.  First imaged by the {\\em Einstein} IPC, the X-ray emission from A2256 was found to depart significantly from spherical symmetry \\citep{FAB84}.  Observations by the R\\\"{o}entgen Satellite (ROSAT) resolved the non-spherical X-ray emission into two maxima in the central region of the cluster \\citep{BRIEL91}. Briel and Henry (1994) used several Position Sensitive Proportional Counter (PSPC) observations with the ROSAT Observatory to obtain a temperature map of A2256.  They found an extremely complicated temperature structure with two opposing 12.0 keV lobes almost perpendicular to a proposed infall direction for the two merging subgroups.  However, the uncertainties of this temperature map are large.  Later observations by the Advanced Satellite for Cosmology and Astrophysics (ASCA) did not observe a similar central temperature structure, but did find a steep radial temperature profile in agreement with ROSAT data \\citep{MARKEVITCH96}.  A2256 also exhibits a number of striking features in the radio regime, as discussed by \\citet{ROTTGERING94}.  It has at least four galaxies with head-tail morphologies, placing it amongst three to four other clusters with comparable numbers of confirmed head-tail radio features.  It also contains one of the longest and narrowest head-tail structures at 700 kpc in length and less than 2.5 kpc in width.  A Z-shaped structure with an ultra-steep spectrum has been tentatively identified with a cluster galaxy, but it eludes definitive classification.  The most pronounced radio features in A2256 are two extended halos with spatial dimensions that are estimated at 1.0 Mpc by 0.3 Mpc.   A2256 is well suited for the study of substructure and the role of dominant galaxies.  Its optical structure is very much like that of the Coma cluster (A1655) with a central pair of dominant galaxies, each at the center of a concentration of fainter galaxies.  The galaxy distribution also indicates the likely presence of the substructure that is seen in the X-ray images.  In this paper, we present new optical imaging of the central region of the cluster and 278 new radial velocities obtained with the WIYN telescope.  In section \\ref{sec:data}, we describe both new imaging and spectroscopy obtained by the Wisconsin-Indiana-Yale-NOAO (WIYN) telescope.  We also include the results of the data analysis.  In section \\ref{sec:analysis}, we review the statistical methods for searching for substructure.  Section \\ref{sec:discussion} includes a discussion of the implications of the observations and the statistical tests.  In section \\ref{sec:summary_and_conclusion}, we summarize the results of this study. ",
        "conclusions": "\\label{sec:summary_and_conclusion}  We have presented the results of a radial velocity survey of the galaxy cluster A2256, carried out with the Hydra multi-object spectrograph on the WIYN telescope, in which we obtained a total of 277 radial velocities of cluster members.  In addition, we presented new photometry and astrometry of 861 cluster members within the central $19\\arcmin \\times 19\\arcmin$ of the cluster, obtained from a $3 \\times 3$ WIYN image mosaic in the R band.  By applying a number of statistical tests, we found strong evidence of statistically significant deviations from the expected single-Gaussian velocity distribution of a relaxed cluster.  On application of the KMM software package, which is based on the EM algorithm for mixture modeling \\citep{MCLACHLAN88}, we found a total of three subclusters, only two of which had been previously noted.  This analysis used both velocity and spatial distribution data to identify the separate groups.  We obtained virial mass estimates for each group and showed that the total of the group masses is less than half the inferred virial mass of the entire cluster, in support of the presence of substructure.  We then applied the two-body orbit model of Beers \\ea\\ (1982), first to the two-body system defined by groups 1 and 2 alone and then to a system in which groups 1 and 2 were treated as a single body and group 3 was taken as the other body.  This dynamical analysis strongly indicates that groups 1 and 2 are an infalling, bound system.  However it does not, in itself, fix the dynamical status of group 3 within the A2256 system.  To help determine whether group 3 is bound to the main cluster, we correlated our optical results with previous X-ray and radio imaging of A2256.  The X-ray surface brightness distribution of A2256 has two centrally located peaks \\citep{BRIEL91}.  The galaxy groups 1 and 2, which we identified from the galaxy position and velocity distributions, are strongly correlated with the X-ray peaks, providing additional support for the merging subcluster interpretation.  The optical data alone allow group 3 to be modeled as either bound, and thus in close proximity with the 1--2 system, or unbound, and thus not interacting with 1--2 system.  However, radio maps of A2256 \\citep{ROTTGERING94} provide strong evidence for an interaction between groups 2 and 3, and thus clearly support the bound model.  As discussed in section \\ref{sub:radio_data}, the diffuse radio emission feature is probably a shock front between the intracluster media of groups 2 and 3.  In addition, the interaction of the head-tail sources C and I with the diffuse radio emission feature provides evidence that the two systems are interacting (see \\S\\ref{subsub:radio_g_h}).  In summary, our proposed model of A2256 consists of three bound subgroups that are undergoing merging.  Groups 1 and 2 appear to be in an advanced stage of the merging process, near the time of the first close passage of the group centers.  In this model, group 1 is located behind group 2 and is infalling from the northwest.  Pinkney \\ea\\ (1996) suggest that the line of centers lies at approximately 45\\arcdeg\\ from the line of sight, which is consistent with the range allowed for bound orbits.  According to our model, group 3 is located on the near side of the 1--2 system and is infalling from the north. Our two-body analysis indicates that the infall direction lies at an angle to the line of sight in the range 50\\arcdeg\\ -- 60\\arcdeg.  The dynamics of the dominant group 2 are probably strongly influenced by the ongoing merger with group 1, but the 1--2 merger system may be largely unaffected by the presence of the much smaller group 3.  Each group contains a central, dominant galaxy which, for groups 1 and 2, is located near the X-ray surface brightness peak. This is consistent with the prediction that the dominant galaxy of a group will quickly settle into the center of the local potential well \\citep{BODE94}.  The greatest offset between the galaxy distribution centroid, the position of the central dominant galaxy, and the X-ray peak occurs in group 1, suggesting that the apparent merger event is strongly influencing the dynamics of this group.  As discussed by \\citet{HENRIKSEN99}, there is some continuing controversy regarding the merger model for interpreting the X-ray observations of A2256. Thus, further X-ray observations of A2256 would help to distinguish between the cases of merging subclusters versus clusters that are merely superimposed on the sky by projection \\citep{HENRIKSEN99}.  Briel \\& Henry (1994) suggest the presence of a complex temperature structure in the X-ray intensity distribution, but this is not strongly supported by observations made by ASCA \\citep{MARKEVITCH96}.  Because of the high velocity impact between groups 1 and 2, the presence of a shock front is likely and has been predicted by computer simulations \\citep{ROETTIGER93}.  A2256 is an excellent target for the Chandra X-ray Observatory.  Higher resolution X-ray images and better temperature maps of the A2256 field would greatly increase our understanding of the dynamics of the multiple merger events that appear to be shaping this system."
    },
    "0201/astro-ph0201160_arXiv.txt": {
        "abstract": "This paper reviews the field of gamma-ray astronomy and describes future experiments and prospects for advances in fundamental physics and high-energy astrophysics through gamma-ray measurements. We concentrate on recent progress in the understanding of active galaxies, and the use of these sources as probes of intergalactic space. We also describe prospects for future experiments in a number of areas of fundamental physics, including: searches for an annihilation line from neutralino dark matter, understanding the energetics of supermassive black holes, using AGNs as cosmological probes of the primordial radiation fields, constraints on quantum gravity, detection of a new spectral component from GRBs, and the prospects for detecting primordial black holes. ",
        "introduction": "With new experiments such as GLAST and VERITAS on the horizon, we are entering an exciting period for gamma-ray astronomy. The gamma-ray waveband has provided a new spectral window on the universe and has already resulted in dramatic progress in our understanding of high energy astrophysical phenomena. At these energies the universe looks quite different then when viewed with more traditional astronomical techniques. The sources of high energy gamma rays are limited to the most extreme places in the universe: the remnants of exploding stars, the nonthermal Nebulae surrounding pulsars, the ultra-relativistic jets emerging from supermassive black holes at the center of active galaxies, and the still mysterious gamma-ray bursters.  While understanding these objects is of intrinsic interest (how does nature accelerate particles to such high energies? how do particles and fields behave in the presence of strong gravitational fields?), these objects can also be used as probes of the radiation fields in the universe and possibly of spacetime itself.  In this case, the astrophysics of the object is a confounding factor that must be understood to produce a quantitative measurement or a robust upper limit.  While some may view this as a limitation of such {\\it indirect} astrophysical measurements, in most cases there are no earth-bound experiments that can probe the fundamental laws of physics at the energy scales available to gamma-ray instruments.  Gamma-ray astronomy has developed along two separate paths. From the ground, simple, inexpensive experiments were built in the 1950's to observe the Cherenkov light generated by extensive air showers generated by photons with energies above several TeV.  Despite decades of effort it was not until the late 1980's that a source of TeV photons was observed.  There are now roughly 10 known sources of TeV gamma rays, three galactic sources and at least three active galaxies. From space, the COS-B satellite, launched in 1975, observed the first sources of cosmic gamma rays at energies above 70~MeV. The launch of the Compton Gamma Ray Observatory (CGRO) in 1991, with the {\\it Energetic Gamma Ray Experiment Telescope} (EGRET) instrument, brought the field to maturity. Whereas COS-B discovered a handful of sources, EGRET observed over 65 active galaxies\\cite{Hartman99}, seven pulsars, many gamma-ray bursts, and over 60 sources that have no known counterparts at other wavelengths.  The disparity in the development of the two techniques can be traced to the extremely low fluxes of particles present above a TeV ($\\sim 4\\gamma$ football field$^{-1}$~hr$^{-1}$) and the cosmic-ray background. Above the earth's atmosphere, one can surround a gamma-ray detector with a veto counter that registers the passage of charged particles. From the ground, one is forced to infer the nature of the primary particle by observing the secondary radiation generated as the extensive air shower develops.  It was not until such a technique was developed for air Cherenkov telescopes~\\cite{weekes-img}, that sources of TeV photons were discovered. Despite these difficulties a new generation of ground-based instruments is under development that will have a sensitivity that will rival that of space-based instruments.  At the same time a space-based instrument, GLAST, with a relatively large area ($\\sim 1$m$^2$) and excellent energy and angular resolution is scheduled to be launched in 2005.  In this paper we will give a brief survey of the gamma-ray universe and demonstrate some of the fundamental measurements (relevant to particle physicists) that can be made using distant objects that emit high-energy photons.  What will hopefully become clear from this exposition are some development paths for future instruments. The need to see to the far reaches of the universe, makes a compelling case for ground-based instruments with energy thresholds as low as 10~GeV. The need to detect and study the many transient phenomena in the universe makes a compelling case for the development of an instrument that can continually monitor the entire overhead sky at energies above $\\sim100$~GeV with sensitivities approaching that of the next generation of pointed instruments.  As with any new branch of astronomy, it is impossible to predict what knowledge will ultimately be gained from studying the universe in a different waveband, but early results hint at a rich future. New and planned instruments with greatly increased sensitivity will allow us to look farther into the universe and deeper into the astrophysical objects that emit gamma rays. Gamma-ray astronomy can be used to study the most extreme environments that exist in the universe, and may also provide a number of unique laboratories for exploring the fundamental laws of physics at energies beyond the reach of earth-bound particle accelerators. ",
        "conclusions": "The field of gamma-ray astronomy has changed radically in the past decade. In 1990 there were a handful of gamma-ray sources observed from space and only one source observed from the ground.  Today there are over 150 known sources of gamma rays observed from space and about 10 sources of TeV gamma rays observed from the ground. The coming decade promises even greater advances. There are new ground-based and space-based instruments currently under development that will increase the number of known sources by over an order of magnitude.  Large field-of-view instruments with relatively low-energy thresholds ($\\sim 500$ GeV) capable of detecting very-high-energy emission from gamma-ray bursts are now beginning to acquire data.  It is clear that this is still a relatively young field were new ideas are being pursued and large advances are on the horizon.  At the workshop we heard from many participants about current measurements and the increasing overlap with particle physics.  The time where distant cosmic sources can be used to probe the fundamental interactions of matter at energy scales unattainable on Earth is approaching. Much discussion centered around possible future directions.  Two goals were put forth for ground-based instruments: a very low energy threshold instrument ($\\sim$10 GeV) and a very sensitive wide-field instrument with an energy threshold $\\sim$100 GeV.  We look forward to the coming decade."
    },
    "0201/astro-ph0201483_arXiv.txt": {
        "abstract": "We present the photometric and spectroscopic evolution of the type IIn \\g\\/ in NGC 1643, on the basis of 4 years of optical and infrared observations.  This supernova shows very flat optical light curves similar to SN 1988Z, with a slow decline rate at all times. The spectra are characterized by strong Balmer lines with multiple components in emission and with a P-Cygni absorption component blueshifted by only 700 \\kms.  This feature indicates the presence of a slowly expanding shell above the SN ejecta as in the case of SNe 1994aj and 1996L.  As in other SNe IIn the slow luminosity decline cannot be explained only with a radioactive energy input and an additional source of energy is required, most likely that produced by the interaction between supernova ejecta and a pre--existent circumstellar medium.  It was estimated that the shell material has a density $n_H >> 10^{8}$ cm$^{-3}$, consistent with the absence of forbidden lines in the spectra.  About 2 years after the burst the low velocity shell is largely overtaken by the SN ejecta and the luminosity drops at a faster rate. ",
        "introduction": "Supernovae of type II (SNe II) are characterized by H lines in their spectra.  Among them a sub-class has been isolated which shows slow photometric evolution, absence of broad P-Cyg absorptions, and emission lines (\\Ha, in particular) with multiple components. Schlegel (1990) labeled these SNe as type IIn, where the ``n'' stands for``narrow'', though this designation is somewhat misleading because of the simultaneous presence of narrow and broad components. Best studied cases of this class are \\z\\/ (Stathakis \\& Sadler, 1991; Turatto et al., 1993; Filippenko, 1997; Aretxaga et al., 1999) and SN~1995N (Fox et al., 2000; Fransson et al., 2001). We recall that some very bright SNe IIn with very broad emission components, implying velocities as large as 17000 \\kms, may be associated with GRB, namely SN~1997cy \\cite{tura00} and SN~1999E \\cite{fili99,cap99a}. Another group, sometimes called SNe IId (or SNe IIsw), shows narrow P--Cygni absorptions on top of otherwise normal SNe II spectra (SN~1984E \\cite{dopi84}, SN~1994aj \\cite{bene98}, SN~1996L \\cite{bene99} and SN~1996al \\cite{ben01b}). These features have been explained by the presence of a thick expanding shell above the photosphere.  The present paper contains an extensive study of \\g, an object which sharing properties both of \\z\\/ and SNe IId reinforces the link between these apparently different objects. It is worth to remind that a few other objects sharing some of the features of \\g\\/ have been reported in recent years, i.e. SN~1997ab (Salamanca et al. 1998), SN~1997eg (Salamanca et al. 2001) and SN~1998S (Fassia et al. 2000, 2001).  SN 1995G was discovered by Evans et al. \\shortcite{evan95} on Feb. 23.5 U.T., using the 1-m reflector of the Australian National University in Siding Spring at an apparent visual magnitude of 15.5. McNaught \\& Cass \\shortcite{cass95} measured the SN position as R.A. = $04^h43^m44^s.22$, Dec = $-05^{o}18'53''.8$ (equinox 2000.0), 4''.5 E and 16''.1 N from the center of NGC 1643 (see Fig. \\ref{paperoga6}), a Sbc galaxy, with $v_{helio}$= 4850$\\pm$29 km/s \\cite{huch93}.  Lacking other determinations, in the following we will adopt the galaxy distance modulus derived from the recession velocity, after correction for the Local Group infall into the Virgo Cluster (LEDA\\footnote{http://leda.univ-lyon1.fr}).  Assuming $H_{0}$=65 km s$^{-1}$ Mpc$^{-1}$, we obtain $\\mu$=34.32 (cfr. Sect.~4.2).  The basic information about \\g\\/ and the parent galaxy is summarized in Tab. \\ref{moe}.  It is worth mentioning that another type II supernova, SN~1999et (Cappellaro, 1999b), has been discovered in the same host galaxy.  We checked that no $\\gamma$--ray burst has been detected by BATSE in the months before the discovery within 2 error radii from the position of \\g. The closest event is burst 4B 950206B, which occurred 18 days before \\g \\/discovery \\/at 2.33 BATSE error radii.  The observations of SN 1995G, described in Sect.~2, are analyzed in Sect.~3 and 4. The discussion of the data is given in Sect.~5, and the conclusions are summarized in Sect.~6.  \\begin{figure}  \\psfig{file=Map_95Gmar95.ps,width=8cm} \\caption{Identification of SN 1995G and the stars of the local sequence around NGC1643 (see Tab.\\protect\\ref{zzz}). North is up, East is to the left.} \\label{paperoga6} \\end{figure}  \\begin{table} \\caption{Main data on SN 1995G and the host galaxy.} \\label{moe} \\begin{tabular}{|c|c|c|} \\multicolumn{3}{|c|}{SN 1995G} \\\\ \\hline $\\alpha$ (J2000.0) & 04h43m44\\fs22 & $\\diamond$ \\\\ $\\delta$ (J2000.0) & $-05$\\degr18\\arcmin53\\farcs8 & $\\diamond$ \\\\ Offset SN - Gal. Nucleus & 4''.5E, 16''.1N & $\\diamond$ \\\\ SN Type & IIn & $\\otimes$ \\\\ Discovery Date & 1995 Feb 23.5 & $\\odot$ \\\\ (Julian Date) & (2449772) & $\\odot$ \\\\ Discovery Magnitude & 15.5 & $\\odot$ \\\\ &&\\\\ \\multicolumn{3}{|c|}{NGC 1643} \\\\ \\hline $\\alpha$ (J2000.0) & 04h43m44s & $\\star$\\\\ $\\delta$ (J2000.0) & $-05$\\degr19\\arcmin10\\arcsec & $\\star$\\\\ Morph. Type & SB(r)bc pec? & $\\star$\\\\ Magnitude & 14.00 & $\\triangle$ \\\\ Galactic Extinction $A_{B}$ & 0.19 & $\\star$\\\\ Diameters & 1\\farcm0 $\\times$ 1\\farcm0 & $\\triangle$ \\\\ $v_{helio}$ (km s$^{-1}$) & 4850$\\pm$29 & $\\triangle$ \\\\ $\\mu$ ($H_{0}$=65 km s$^{-1}$ Mpc$^{-1}$) & 34.32 & $\\dag$\\\\ \\hline \\end{tabular}  $\\diamond$ \\protect\\cite{cass95}\\\\ $\\otimes$ \\protect\\cite{fili95}\\\\ $\\star$ \\protect\\cite{deva91}\\\\ $\\odot$ \\protect\\cite{evan95}\\\\ $\\triangle$ \\protect\\cite{huch93}\\\\ $\\dag$ LEDA\\\\  \\end{table}  \\begin{figure} \\psfig{file=SN95G_301196_IR2.eps,width=7.6cm} \\caption{SN 1995G and the IR local sequence (J band image, taken with ESO 2.2m + IRAC2 on 1996 November 30).  Note that star A coincides with star 1 of the optical sequence. IR magnitudes of the stars are given in Tab.\\protect\\ref{azz}.} \\label{willie} \\end{figure} ",
        "conclusions": "This paper presents the photometric and spectroscopic observations of \\g\\/ obtained at ESO La Silla, Lick and Asiago, over a period of more than 4 years after the discovery.  The broad band and bolometric light curves are flatter than those expected from radioactive decay. The luminosity peak and evolution are very similar to those of the well studied \\z. This implies the presence of an additional source of energy, most likely the interaction of the ejecta with a dense CSM.  The line profiles show evidence of different components. The strongest line, \\Ha, shows relatively narrow P-Cygni profiles with minima displaced by 700-800 \\kms superimposed on an intermediate width emission FWHM$ \\approx2000$ \\kms and on a broader component having FWHM$ \\approx4000$ \\kms. The widths of these emission lines decrease with time.  At 942 days after maximum the \\Ha emission has a boxy shape with FWZI=2500 \\kms, suggesting that most of the expanding shell has been swept away and the optical depth of the material is low.  The emission lines of \\g\\/ were narrower than in \\z, indicating slower expansion velocities of the ejecta.  The detection of strong He lines between about 200 and 600 days after the discovery indicates the presence of non-thermal effects. Simple considerations about the density of the slowly expanding material indicate densities $n_H >> 10 \\time 10^8$ cm$^{-3}$, in agreement with the absence of forbidden lines.  In the hypothesis that the additional source of luminosity of \\g\\/ is the conversion to radiation of the kinetic energy of the ejecta due to interaction with a precursor wind, we estimate the mass of the ejecta and the mass loss. We obtain $M_{ej} \\simeq 3-9$ $M_{\\odot}$ and $\\dot{M}$ = 0.002 $M_{\\odot}$/yr, respectively, provided that the conversion efficiency is $\\psi=1$ and the explosion energy is $E = 10^{51}$erg."
    },
    "0201/astro-ph0201246_arXiv.txt": {
        "abstract": "In October 1997, the Italian and Dutch GRB teams started a collaboration on ESO optical follow-up of rapidly and accurately localized GRBs. Subsequently, starting April 1, 2000, this collaboration was extended to astronomers from other countries, who contributed their expertise for the creation of a Consortium committed to the study of GRB counterparts and host galaxies at optical and near-infrared wavelengths. The collaboration aims at the joint exploitation of the observations taken within an ESO Large Programme approved for the two-year period April 1, 2000 - March 31, 2002. Here we describe history and organization of this Consortium, the goals of the ESO Large Programme, and the main results obtained up to now with ESO telescopes. ",
        "introduction": "\\vspace{-.2cm}  The availability of fast (few hours after the GRB trigger) and precise (arcmin-sized) GRB localizations, first afforded by the Italian-Dutch X--ray satellite {\\it BeppoSAX} in 1996 and subsequently by other spacecraft, allowed astronomers to pinpoint Optical Transients (OTs) associated with the high energy events and to effectively explore the physics behind these phenomena. In this contribution we briefly outline history, status and results of the search and followup of GRB optical and near-infrared (NIR) counterparts at ESO since the beginning of the {\\it BeppoSAX} afterglow era.  The first ESO ToO observation of a fast and precise GRB position was activated on January 1997 at NTT on the GRB970111 error box, but no optical counterpart was detected. The first detection of an OT associated with a GRB was achieved at ESO about two months later, on 1997 March 13, with NTT observations of the GRB970228 error box \\cite{jvp}, the first GRB for which X--ray and optical afterglows were discovered.  Subsequently, thrusted by this outstanding scientific achievement, the Italian and Dutch GRB search and followup teams, led by Filippo Frontera and Jan van Paradijs, respectively, independently submitted regular ESO proposals for Period 60 (October 1997 - March 1998) for the activation of ToO observations at several ESO telescopes on GRB error boxes observable from Chile. Both proposals were approved and, following the suggestion from ESO, the Italian and Dutch groups decided to start collaborating in this search and to form a single team.  The Italian-Dutch collaboration in the GRB follow-up at the ESO telescopes was organized in a way that the two groups alternated in the program lead at every trigger of an accurately and rapidly localized GRB, independent of spacecraft. The leadership of the group `on duty' encompassed every step from alerting the telescopes to taking responsibility of results publication.  Besides the Italian-Dutch collaboration, a parallel proposal set up by Danish, Spanish and German teams was active at several ESO telescopes.  During Periods 60 to 64 (October 1997 - March 2000) ESO observations placed several milestones in the study of GRB optical afterglows, in particular: the detection of SN 1998bw in the error box of GRB980425, the only case in which a SN was clearly detected in the field of a GRB within stringent temporal and spatial limits \\cite{gal}; the discovery and monitoring of optical polarization in the OTs of GRB990510 \\cite{wij} and GRB990712 \\cite{rol}; the determination of the redshift of GRBs 990510, 990712 \\cite{pmv00} and 991216 \\cite{pmv01} using VLT-Antu; the NTT detection of two very red afterglows associated with GRB980329 \\cite{pal} and GRB990705 \\cite{mas}, their color being most likely due to high local absorption in the host galaxy or to high redshift; and the discovery of the farthest GRB observed so far, located at redshift $z$ = 4.5 \\cite{and}. Also, in several cases the detection and the spectrophotometric observations of host galaxies associated with GRBs were accomplished with VLT-Antu and NTT ESO telescopes.  \\vspace{-.2cm} ",
        "conclusions": ""
    },
    "0201/astro-ph0201300_arXiv.txt": {
        "abstract": "We have studied the long-term X-ray light curve (2$-$10 keV) of the luminous Seyfert 1 galaxy MCG-2-58-22 by compiling data, from various X-ray satellites, which together cover more than 20 years. We have found two distinct types of time variations in the light curve. One is a gradual and secular decrease of the X-ray flux, and the other is the episodic increase of X-ray flux (or flare) by a factor of 2$-$4 compared with the level expected from the secular variation. We detected 3 such flares in total; a representative duration for the flares is $\\sim $2 years, with intervening quiescent intervals lasting $\\sim\\! 6-8$ years. We discuss a few possible origins for these variabilities. Though a standard disk instability theory may explain the displayed time variability in the X-ray light curve, the subsequent accretions of stellar debris, from a tidal disruption event caused by a supermassive black hole in MCG-2-58-22, cannot be ruled out as an alternative explanation. ",
        "introduction": "X-ray observations of active galactic nuclei (AGN) show that many of them are variable, over a range of amplitudes, and on many timescales (Lawrence et al.\\ 1985; Grandi et al.\\ 1992; Mushotzky, Done, \\& Pounds 1993; Nandra et al.\\ 1997; Ulrich, Marachi, \\& Urry 1997; Ptak et al.\\ 1998; Turner et al.\\ 1999). Variability of the X-ray emission is a powerful probe of physical processes occurring in the inner regions of AGN. In particular, rapid variability is widely thought to be related to the central regions, and it has actually been used to constrain the physical properties of the central engine. For example, the short-time variability amplitudes of Seyfert 1 galaxies are known to be anti-correlated with the source luminosities (Barr \\& Mushotzky 1986; Nandra et al.\\ 1997; Ptak et al.\\ 1998). This correlation may reflect differences in the masses of the central supermassive black holes (SMBHs). In addition to the short-time variabilities, the study of long-time variabilities is also important since it may bring out other interesting information, such as the global structure of the accretion disk around the central SMBH, or episodic events such as flares and outbursts.  Detailed studies of the long-term X-ray light curves of AGN have begun only recently. Markowitz \\& Edelson (2001) analyzed 300-day light curves of Seyfert 1 galaxies in the 2$-$10~keV band. They showed that the X-ray variability of Seyfert 1 galaxies is described by a single, universal power-density spectrum (PDS), and that the cutoff moves to longer timescales for sources with higher luminosity. Soft X-ray outbursts, having an amplitude of about 2 orders of magnitude, were observed from NGC 5905 (Bade, Komossa, \\& Dahlem 1996) and Zwicky 159.034 (Brandt, Pounds, \\& Fink 1995). Various possible scenarios for such outbursts have been reviewed by Komossa \\& Bade (1999). However, long-term modulations in AGN light curves, which exceed several years, have mostly been studied in the optical range. For instance, Webb (1990) reported the results of 61~years of optical observations of 3C 120, and claimed the presence of three variability components: a sinusoidal component with a period of 12.43 years, a linear component, and high amplitude flares on much shorter timescales. Peterson et al.\\ (1998) reported the spectroscopic monitoring of nine Seyfert 1 galaxies in the optical band, the aim of which was to determine the size of the broad line emission regions. In order to perform a similar analysis in the X-ray band, we need to collect all available datasets from the various X-ray satellites. Because the different instruments of the satellites have their own individual properties, and cover different energy ranges, careful analysis is required for such studies in the X-ray region.  In this paper, we study the long-term X-ray light curve of MCG-2-58-22, covering more than 20 years, and we discuss the potential origins of the long-term behaviors. In previous work, Choi et al.\\ (2001) analyzed the Ginga, ROSAT, and ASCA data for this source, and noticed flare-like events. This motivated us to perform a thorough analysis of the long-term X-ray flux variations in this source. For this purpose, we have gathered X-ray flux measurements of MCG-2-58-22 from the literature, as well as raw X-ray data from the HEASARC public archives at NASA/GSFC, and from the SIRIUS database at ISAS\\@. The observational data we gathered include those of HEAO-1 (Rothschild et al.\\ 1979), Einstein (Giacconi et al.\\ 1979), EXOSAT (Taylor et al.\\ 1981), Ginga (Turner et al.\\ 1989), ROSAT (Pfeffermann 1986), ASCA (Tanaka, Inoue, \\& Holt 1994), and RXTE ({\\tt http://heasarc.gsfc.nasa.gov/docs/xte}).  MCG-2-58-22 is a luminous Seyfert 1 galaxy at $z=0.04732\\pm0.0003$ (e.g. Huchra et al.\\ 1993). The X-ray luminosity of MCG-2-58-22, $L_X \\sim 10^{44}$ erg s$^{-1}$, is known to be variable on a timescale of a few times $10^{3}$ seconds to years (Grandi et al.\\ 1992; Nandra \\& Pounds 1994; Choi et al.\\ 2001). Many observational characteristics are typical of Seyfert 1 galaxies: time variability, a power-law type continuum spectrum, and a soft excess phenomenon (Ghosh \\& Soundararajaperumal 1992; Nandra \\& Pounds 1994; Weaver et al.\\ 1995; Nandra et al.\\ 1997; George et al.\\ 1998; Turner et al.\\ 1999). The mass of the putative central SMBH of MCG-2-58-22 is estimated from the UV and optical observations to be a few times $10^8$~M$_{\\odot}$ (e.g., Padovani \\& Rafanelli 1988; Wandel 1991). MCG-2-58-22 has also been studied in other wavelength bands. Mundell et al.\\ (2000) observed MCG-2-58-22 using the VLBA at 8.4~GHz, and detected the nuclear radio source, which was without any extended structures, in the parsec-scale image; this suggests that the VLBA radio source could be a ``bare'' Seyfert 1 nucleus (see also Weaver et al.\\ 1995). In the optical region, this source has displayed a continuum variation on a timescale of $\\sim1$ year (de Ruiter \\& Lub 1986; Whittle 1992), as well as a very wide variability in the Balmer line profiles and luminosities (see, e.g., Winkler et al. 1992, and references therein). ",
        "conclusions": "As shown in Figure~1, MCG-2-58-22 clearly shows two characteristic variabilities: the gradual, secular decrease of the X-ray flux, and multiple flares with a representative duration of $\\sim$2~years. These two characteristic variabilities imply that the source undergoes at least two distinctly different physical processes. Magnetic reconnection may be considered as a mechanism for the flares, but it seems unlikely that this process could explain their duration and repetition frequency, since the magnetic field evolution timescale in an accretion disk is of order the dynamical time (Romanova et al. 1998; Poutanen \\& Fabian 1999). Another possible, and more likely, origin of the variabilities are the instabilities which could arise from the accretion disk. Such instabilities may result in the modulation of the mass-accretion rate, leading then to the observed flux variations.  Disk instabilities which could be appropriate for the variabilities in MCG-2-58-22 are the viscous-thermal and the viscous instabilities. If the disk experiences a viscous-thermal instability caused by a sudden change in the hydrogen ionization state, we estimate the instability timescale to be $t_{\\rm vis-th}\\sim 6\\times 10^2(\\alpha/0.1)^{-1} M_8^{1/3} ({\\dot M}/10^{-4}M_{\\odot}\\ {\\rm yr}^{-1})^{1/3}$~yr, where $M_8$ is the SMBH mass in units of $10^8$~M$_{\\odot}$. On the other hand, the viscous timescale near the innermost radius of the standard $\\alpha$-disk can be a few tens of years for a black-hole mass of $10^8$~M$_\\odot$. Because these timescales are much longer than the observed duration of the flares, the viscous-thermal and the viscous instabilities are probably not the cause of the flares. However, the gradual and secular decrease of the X-ray flux has a longer timescale, and this slower variation could be caused by these types of disk instability.  Temporal increases of the mass accretion rate could arise, in principle, from other mechanisms. One entertaining possibility is the tidal disruption of stars by the SMBH (e.g., Rees 1988; Lee \\& Kim 1996; Kim et al.\\ 1999). In this model, the frequency with which a star passes within a distance $r$ from the SMBH can be estimated to be $\\approx 10^{-3}~M_8^{4/3} (N_*/10^6~ {\\rm pc^{-3}}) (\\sigma/300~ {\\rm km~s^{-1}}) (r/r_t)~~ {\\rm yr^{-1}}$, where $N_*$ is the number density of the stars, $\\sigma$ is the virial velocity of the stars, and $r_t$ is the tidal radius of the SMBH\\@. Although the number density of the stars and their velocity dispersion near the SMBH are not known, the flare events we observed from MCG-2-58-22 may be difficult to interpret as independent tidal disruption events, because the event rate would then be too high. Instead, the resulting flares may be produced when the bound material from the tidally disrupted star returns to the pericenter, with an orbital period of a few years (Rees 1988; Ulmer 1999). The observed peak luminosity of $\\sim\\! 4 \\times 10^{44} ~{\\rm erg~s^{-1}}$ for the flares corresponds to that from the debris of $\\sim$0.1~M$_\\odot$ being swallowed steadily with 10~\\% efficiency over a year's duration. However, it is known that a star may be captured by the SMBH without tidal disruption, provided the SMBH is a Schwarzschild black hole heavier than $\\sim\\!10^8$~M$_\\odot$; this limit is comparable to previous estimates of the SMBH mass in MCG-2-58-22. If we consider a Kerr black hole, the tidal disruption may still be possible, depending on the trajectory of the approaching star (Beloborodov et al.\\ 1992). Moreover, the atmospheres of giant stars could be stripped off before being captured by the SMBH, and this process may be the actual cause of the flares (e.g., Nolthenius \\& Katz 1982; Carter \\& Luminet 1983; Rees 1988; Laguna et al.\\ 1993; Ulmer 1999). Thus, the observed flares are not inconsistent with those caused by the tidally disrupted stellar debris near the SMBH, even though the flare properties would depend on many unknown parameters, such as: the type of the disrupted star, the spin of the SMBH, and the minimum radius of the trajectory.  The long-term optical light curve of MCG-2-58-22 was obtained by Winkler et al.\\ (1992), as a result of a 4-year campaign to monitor 35 southern Seyfert galaxies. The campaign period from 1987 through 1990 overlaps with the Ginga observations of 1989 and 1990. During this time, a clear variation of about 0.3 mag is reported in the V-band. Unfortunately, however, the optical light curve does not cover the time periods of the X-ray flares we detect, and it shows only gradual and smooth variations. Over this time, the X-ray flux varied by about 40~\\%. This is comparable to the 0.3 mag variations in the V-band. Thus, at least the fractional amplitude of the long-term variations may be similar between the optical and X-ray wavelength bands. It is difficult to ascertain whether the optical variation is physically related to that in X-rays, since the optical observation period is too short, compared with that of the X-ray light curve.  The variabilities we detected from MCG-2-58-22 may be reminiscent of the rapid time variations seen in the galactic black hole candidates (GBHCs). Phenomenologically, the presence of a characteristic timescale is known both in GBHCs in the hard state, and in the SMBHs (Hayashida et al.\\ 1998; Markowitz \\& Edelson 2001). In both kinds of sources, the characteristic timescale corresponds to the cutoff in the PDS, and this timescales may be scaled almost linearly with the black hole mass. The canonical galactic black hole, Cyg X-1, contains about 10~M$_\\odot$, and it shows a PDS cutoff at around $\\sim$10~sec. The PDS of Cyg X-1 may be reproduced by the superposition of shot noise impulses, whose typical duration determines the PDS cutoff frequency (e.g.\\ Negoro, Kitamoto, \\& Mineshige 2001). If we assume that the SMBH in MCG-2-58-22 has a mass of $10^8$~M$_\\odot$ (e.g., Padovani \\& Rafanelli 1998; Wandel 1991), the characteristic timescale is about a few years. This is just about equal to the duration of the flares we detected. Thus the flares in MCG-2-58-22 may be analogous to the shot noise seen in galactic black holes in their hard state.  In conclusion, we have detected two characteristic variabilities in the long-term X-ray light curve of MCG-2-58-22, by analyzing archival data from various X-ray satellites. One variation is the gradual, secular decrease of the X-ray flux, which may have a timescale of several decades, and the other is the flaring. Although it is difficult to accurately measure the duration of, and intervals between, the flares, due to the sparse sampling in the observational data, a representative duration is $\\lesssim $2 years, with an intervening interval of $\\lesssim\\! 6-8$ years. These two distinct timescales may be accounted for by a model with a supermassive black hole accompanied by an unstable accretion disk; the long-term secular variation would be expected from instabilities in the disk, while the short-term flaring would arise from the tidal disruption of stars by the supermassive black hole. Further observations and spectral analysis in other wavelength bands should be explored, in order to verify this scenario."
    },
    "0201/astro-ph0201136_arXiv.txt": {
        "abstract": "The exploration of the end of the Dark Ages will be one of the most exciting field of the next decade.  While most of the proposed observations must await the next--generation telescopes, the observational window of the redshifted 21cm line offers the possibility to investigate the physics of reheating and reionization on a short term.  Here we describe several possible signatures detectable in the wavelength range 100-200 MHz.  Among the physics that can be investigated: the epoch of reheating and reionization; topology and timescales of reheating; the nature of the ionizing sources; the baryon distribution at redshift $z\\sim 10$.  Such a good deal of information is within reach of present--day, or near--future radio facilities. ",
        "introduction": "The Dark Ages are ended by the appearance of the first stars and/or quasars, that reheat the diffuse cosmic baryons and then reionize the Universe (see Rees 1999).  The reionization is defined as the epoch when the volume--filling factor of HII regions is $\\simeq 1$, and the Universe becomes transparent to ionizing radiation (for an extensive review see Loeb \\& Barkana 2001).  The absence of a Gunn--Peterson (GP) effect in the line of sight of distant quasars, put the epoch of reionization ($z_{reion}$) at redshifts $>5$.  To date, there are little additional constraints on the physics of the cosmic baryons at such high redshifts.  This is also due to the complexities involved by any theoretical model that must include the nature and the birthrate of the first luminous objects, their spectrum and emissivity, their feedback into the surrounding Intergalactic Medium (IGM), etc.  A significant step forward has been recently obtained with Keck spectroscopy of the most distant SLOAN quasars (Becker et al. 2001; Djorgovski et al. 2001).  The presence of a sudden increase in the Ly$\\alpha$ opacity between redshift 5 and 6 may indicate that the Universe is approaching reionization at $z\\sim 6$.  However, Barkana (2001) pointed out that, due to the high Ly$\\alpha$ opacity, these observations are also consistent with a post--reionization phase, and that a conclusive proof requires the observation of similar GP absorptions along several additional line of sights.  Alternatively, the observation of smoking--gun signatures has been proposed.  Among them: dectection of scattered Ly$\\alpha$ emission around sources beyond $z_{reion}$ (Loeb \\& Rybicki 1999); features in galaxy number counts at $z>z_{reion}$ (Barkana \\& Loeb 2000).  These observations would be effective probes of the physics of cosmic baryons and of the nature of the first ionizing sources at the same time, but they must wait for NGST and/or SIRTF, which will be operating at the end of this decade.  Signatures of the reionization are expected also in the CMB. MAP and Planck can probe the optical thickness at $z_{reion}$ in the 30--150 GHz range.  In particular, fluctuations on $\\le 0.1$ degree reflect the topology of reionization.  However, such measures are highly degenerate with cosmological parameters.  It is clear that the investigation of the Universe at $z\\simeq 10$ is a difficult task with present--day facilities.  In this situation, the observational window offered by the redshifted 21cm line can open much of the Universe to a direct study of the reheating and reionization epochs.  In the following we will discuss some observations of HI at high $z$ that can shed light on many aspects of these processes. ",
        "conclusions": "The reionization is a crucial epoch in the evolution of cosmic structures.  The same ingredients that characterize it (for example: nature of the first sources, the role of black holes in galaxy formation, the role of feedback, the chemical enrichment of the IGM) will deeply affect the evolution of the Universe in the following epochs. Getting direct information about these crucial early stages of the luminous Universe is an obvious goal for the next--decade cosmology.  To date, several strategies have been proposed, but most of them must wait for the next generation of astronomical facilities.  The observations of high--$z$ HI in the redshifted 21cm, on the other hand, offer the possibility to investigate the reionization epoch using future but also present--day radio facilities.  All--sky signals from HI at redshift $z\\simeq 10$ are recognizable as discontinuities superimposed on the CMB in the wavelength range $\\simeq 100-200$ MHz. In particular, the epoch of reheating can be seen as a deep ($\\simeq -40$ mK) absorption feature against the CMB, at the corresponding redshifted 21cm line.  The density perturbation field at a redshift $z\\approx 5\\div 20$ can be reconstructed looking for mK fluctuations at $1'$--$5'$ resolution in the radio sky.  Finally, luminous quasars can be seen by identifying peculiar, ring--shaped signals whose morphology depends on the source's age, luminosity and geometry.  Such observations represent an exciting scientific and technological challenge, and will constitute a unique investigation of the Dark Ages, for many aspects complementary to future observations at other wavelengths."
    },
    "0201/astro-ph0201185_arXiv.txt": {
        "abstract": "We present archival {\\it Hubble Space Telescope} images of the nuclear regions of 43 of the 46 Seyfert galaxies found in the volume-limited, spectroscopically complete CfA Redshift Survey sample.  Using an improved method of image contrast enhancement, we create detailed high-quality ``structure maps'' that allow us to study the distributions of dust, star clusters, and emission-line gas in the circumnuclear regions (100-1000 pc scales) and in the associated host galaxy. Essentially all of these Seyfert galaxies have circumnuclear dust structures with morphologies ranging from grand-design two-armed spirals to chaotic dusty disks.  In most Seyferts there is a clear physical connection between the nuclear dust spirals on hundreds of parsec scales and large-scale bars and spiral arms in the host galaxies proper.  These connections are particularly striking in the interacting and barred galaxies.  Such structures are predicted by numerical simulations of gas flows in barred and interacting galaxies, and may be related to the fueling of AGN by matter inflow from the host galaxy disks.  We see no significant differences in the circumnuclear dust morphologies of Seyfert 1s and 2s, and very few Seyfert 2 nuclei are obscured by large-scale dust structures in the host galaxies.  If Seyfert 2s are obscured Seyfert 1s, then the obscuration must occur on smaller scales than those probed by HST. ",
        "introduction": "\\label{sec:intro}  There is now a large and compelling body of observational evidence that suggests that most, if not all, galaxies contain supermassive black holes at their centers \\citep[e.g.,][]{richstone98}.  However, active galactic nuclei (AGN) are only found in a small minority of all galaxies in the local universe \\citep[e.g.,][]{huchra92,ho97}.  What is it that makes some galaxies AGN but most others quiescent?  One line of inquiry is to ask if the differences are to be found in their circumnuclear environments.  In particular, is the difference simply a matter of whether or not the central black hole is being provided with interstellar gas to fuel the nuclear activity?  The problem of providing fuel to an AGN from the vast reservoirs of interstellar gas found in the disks of spiral galaxies is how to remove the angular momentum from the gas so it can fall into the nucleus.  The two classical mechanisms that are invoked are interactions \\citep{toomre72} and bars \\citep{schwartz81}, including nuclear bars \\citep{shlosman89,pfenniger90}.  There is an extensive literature devoted to demonstrating that both are theoretically viable mechanisms for fueling AGN \\citep{hernquist89,barnes91,athanassoula92,friedli93,piner95,hernquist95}. However, neither interactions nor bars of either type are sufficiently common among AGN compared to non-AGN galaxies to be the fueling mechanism in all cases \\citep{adams77,petrosian82,keel85,fwstocke88,kotilainen92,mcleod95,keel96, alonso96,mulchaey97,derobertis98,regan99,knapen00,schmitt01,laine01}.  Previous investigations have used {\\it HST} to study the circumnuclear environments of Seyfert galaxies and search for differences between the Seyferts with and without broad-line components. \\citet{nelson96} obtained pre-COSTAR imaging of a large sample of Seyfert and non-Seyfert Markarian galaxies to look for differences in the nuclear structure of these galaxies at higher angular resolution than is possible with ground-based imaging. They discovered that the nuclei of broad-line Seyfert $1-1.5$ galaxies are dominated by strong point sources. In contrast, Seyfert $2$ galaxies and other Markarian galaxies that lack broad-line regions contained weak or no strong nuclear source superimposed on the underlying galaxy's surface brightness profile. This result is further borne out in the extensive {\\it HST} snapshot program of \\citet{malkan98}. They also invariably find more strong central point sources in Seyfert 1 galaxies than in the Seyfert 2s.  As this survey was carried out with the unaberrated WFPC2 PC camera, these investigators were also able to look for differences in the nuclear environments of these galaxies. They found Seyfert 2 galaxies were more likely to possess dusty nuclear environments than Seyfert 1 galaxies, lending support to unified models, which propose obscuration of the broad-line region in Seyfert 2s by dust. Their observations are evidence for the presence of dust on large scales in the nuclear region, not in a torus immediately outside the broad-line region.  Visible--near-infrared color maps obtained with the {\\it Hubble Space Telescope} ({\\it HST}) have shown that the circumnuclear ($\\sim 100 - 1000$ pc) regions of a large number of low-luminosity AGN contain nuclear spiral dust lanes. These spirals are distinct from the spiral arms on kpc scales in the main galaxy disk. $V-H$ color maps of these galaxies show that these `nuclear spirals' extend from 100's of pc scales into the unresolved nucleus \\citep{quillen99,regan99}. Theoretical models for the formation of nuclear spiral structure suggest that it is dynamically distinct from the main disk spiral arms \\citep{bertin89,elmegreen92}. \\citet{martini99} showed that nuclear spirals in AGN reside in nonself-gravitating disks and are therefore likely due to shocks in nuclear gaseous disks. They postulated that as shocks can dissipate energy and angular momentum, these nuclear spirals may be the signature of the fueling mechanism in these galaxies.  Nuclear spirals in AGN have generally been seen mostly in Seyfert 2s, although this could easily be a selection effect: the samples of \\citet{quillen99} and \\citet{regan99} were mostly comprised of Seyfert 2s. Also, \\citet{martini99} only observed Seyfert 2s with {\\it NICMOS} as they have fainter nuclear PSFs, and thus the circumnuclear environments of Seyfert 2s are easier to study with {\\it HST} than the circumnuclear environments of Seyfert 1s. The question remains, however, if Seyfert 1s and 2s contain nuclear spirals with the same relative frequency; that is, nearly 100\\% as seen in Seyfert 2s. As we only have near-infrared {\\it NICMOS} imaging of the Seyfert 2s in the CfA sample, we need to use an alternate technique to look for nuclear spiral structure in these Seyferts.  In \\S\\ref{sec:procedure} we discuss our data-processing, and introduce a technique for creating ``structure maps'' in \\S\\ref{sec:stmaps} that are an excellent surrogate for color maps for detecting small-scale dust-extinction and emission-line features present in the visible-band images.  We then use this technique to compare the circumnuclear environments of the Seyfert 1s and 2s in \\S\\S\\ref{sec:morph} and \\ref{sec:sey}, connecting the nuclear structures seen to the larger host galaxies in most cases. In \\S\\ref{sec:conc} we present a summary of our results, and discuss the implications for the fueling of the active nuclei. ",
        "conclusions": "\\label{sec:conc}  Using a new ``structure mapping'' technique, we have found nuclear spiral dust structures that are plausibly related to the inflow of interstellar gas from the host galaxies into the nuclear regions of a well-defined sample of Seyfert galaxies.  We find the dust morphology that is expected if interstellar gas is being driven primarily by large-scale bars, and if interstellar gas is mass-transported inwards by torques arising from tidal interactions.  There is also circumstantial evidence for the spiral infall of gas, with greater or lesser degrees of coherence, in galaxies which show neither stellar bars nor evidence of tidal interaction.  In most cases the spiral dust structures seen previously at small scales in more traditional color maps are shown to be connected to the large-scale properties of the galaxies, although they are not simply continuations of the large-scale features.  All of these structures are consistent with the idea that interstellar gas from the host galaxy may be transported into the nucleus, to varying degrees of intensity and coherence, via a variety of mechanisms.  The case that these structures are related to fueling the AGN is primarily phenomenological.  Theoretical models of gas inflow in barred and interacting galaxies predict structures similar in appearance to those we see in our images.  However, what we do not know is whether we can make the final connection between these larger-scale structures and the central black hole.  That we find these structures in most of the Seyferts studied here makes this connection plausible, but we cannot yet make it conclusively.  From an observational perspective, a particularly difficult challenge is how to establish kinematic evidence of inflow in cold interstellar material on these scales.  From the theoretical side, work on what inflow rates are expected, and what observational factors (either in addition to or in the absence of direct kinematical measurements) might help inform such estimations on an object-by-object basis.  It has long been known that bars and interactions occur in inactive galaxies, and there is no reason to expect that they will not funnel gas into the inner regions of those galaxies in much the same way.  Further, nuclear spirals have been observed in a number of normal spiral galaxies without obvious AGN activity \\citep{phillips96,carollo98,elmegreen98,laine99}.  If nuclear spirals are signatures of shocks that can dissipate sufficient angular momentum to fuel a black hole, the question remains why normal galaxies that exhibit such structures are not active, particularly as it is now clear that essentially all galaxies harbor nuclear black holes, with masses closely correlated with their host galaxy properties \\citep[e.g.][]{richstone98}.  This suggests that the transport of gas into the central ($<100$pc) region is a necessary, but not unique requirement of nuclear activity.  The near-ubiquity of coherent, circumnuclear dust structures suggestive of shocks and matter inflow supports the conjecture that they are responsible for fueling their AGN. The first step to test this is a study of a well-defined control sample of quiescent spiral galaxies of similar morphological type observed and analyzed in similar ways to the Seyferts.  We are currently engaged in an on-going HST snapshot program to address this question of the relative frequency and strength of nuclear dust spirals in normal galaxies compared to AGN.  More data will be forthcoming."
    },
    "0201/astro-ph0201466_arXiv.txt": {
        "abstract": "We have carried out numerical simulations of freely decaying magnetohydrodynamic (MHD) turbulence in three dimensions, which can be applied to the evolution of stochastic magnetic fields in the early Universe. For helical magnetic fields an inverse cascade effect is observed in which magnetic helicity and energy is transfered from smaller scales to larger scales, accompanied by power law growth in the characteristic length scale of the magnetic field. The magnetic field quickly reaches a scaling regime with self-similar evolution, and power law behaviour at high wavenumbers. We also find power law decay in the  magnetic and kinematic energies. ",
        "introduction": "Magnetic fields are ubiquitous in the Universe, being observed in nature on scales from planetary size to galaxy cluster size \\cite{Zel88,Kro94}. In galaxies and galaxy clusters, the typical strength is of order a few $\\mu$Gauss, which is thought to be produced by dynamo action on a seed field. In galaxies the dynamo timescale is roughly a rotation period, $10^8$ Yr, and a simple calculation \\cite{Zel88} based on the age of a typical galaxy shows that the seed field must have been about \\(10^{-20}\\) Gauss, or perhaps less in the currently favoured models with a cosmological term \\cite{Davis:1999bt}.  There is no shortage of ideas for generating this seed field.  The more conventional astrophysical explanations are based on a Biermann battery operating at the era of reionisation (see e.g.\\ \\cite{GneFerZwe00} and the references therein).  There are more speculative ideas based on various models of inflation \\cite{InfMag}, phase transitions \\cite{PhaTraMag} and primordial black holes \\cite{BlaHolMag}.  All these mechanisms have the common feature of producing stochastic, homogeneous and isotropic magnetic and velocity fields which can be characterised by their power spectra and characteristic initial scales. Our interest here is to try and make model-independent statements about the evolution of the magnetic fields once they are generated.  This article, which is based on Ref.~\\cite{ChrHinBra01}, studies the evolution of a stochastic magnetic field generated at a phase transition, such as the confinement transition in QCD at $t \\simeq 1$ sec, or the electroweak symmetry-breaking transition at $t \\simeq 10^{-11}$ sec.  It therefore falls into the category of decaying 3D MHD turbulence, which has been studied before in the MHD community \\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 focus here on the transfer of magnetic energy from small to large scales, as necessitated by the conservation of magnetic helicity. 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}.  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 the coherence scale of the field growing as $t^{0.5}$, and equivocal evidence for a weak cascade when only helicity fluctuations are present. In the helical case we also find a self-similar power spectrum with an approximately $k^{-2.5}$ behaviour at high $k$. We find decay laws for the magnetic and kinetic energies of $t^{-0.7}$ and $t^{-1.1}$, respectively, in the helical case, and $t^{-1.1}$ for both in the non-helical case. ",
        "conclusions": "The most directly comparable simulations of decaying 3D MHD turbulence were carried out by Biskamp and M\\\"uller \\cite{BisMul99,MulBis00}. They found similar results, the magnetic field evolved self-similarly, with a power-law behaviour at high $k$.  However, their power law was $k^{-5/3}$, much less steep than our $k^{-2.5}$.  There were a number of differences between their and our simulations:  they were able to achieve larger Reynolds numbers, both by having larger grids, and by using hyper-diffusivity (a $\\nabla^4$ magnetic diffusivity term). However, we believe that the real difference is due to their initial cut-off scale being significantly larger, at $k_c=4$. We have performed a run with larger initial length scale, $k_{c} = 5$. In this case the magnetic energy spectrum develop into an approximate $k^{-5/3}$ law at late times. However, this occurs only after the peak of the spectrum has left the simulation box.  There have also been renormalisation group (RG) analyses looking for an inverse cascade in driven MHD turbulence \\cite{Shiromizu:1998bc,Berera:2001eb}. In particular Berera and Hochberg \\cite{Berera:2001eb} saw no evidence for an inverse cascade. However, it is not clear that the results are directly comparable, firstly because we are considering freely decaying turbulence, and secondly because RG analysis can only give information about late times when the system is in equilibrium with the driving force. What we have referred to as an inverse cascade is, even in the driven case \\cite{Bra01}, a time-dependent phenomenon characterised by a bump in the power spectrum travelling to smaller wave number.  In conclusion, we find good evidence from our numerical simulations that helical stochastic magnetic fields show an inverse cascade (in the sense explained above), and that even if only small helicity fluctuations are present initially, there is still weak inverse cascade.  We have determined growth laws for the magnetic and kinetic energies $E_M$ and $E_K$. In the helical case, \\(E_M \\sim t^{-0.7}\\) and \\(E_K \\sim t^{-1.1}\\), which means that that there is no equipartition of energy. This is because the extra constraint of helicity conservation forces the magnetic field to transfer power to larger scales rather than allow it to be dissipated.  The importance of helicity is borne out by the fact that in the non-helical case, we find \\(\\xi \\sim t^{0.4}\\) and \\(E_M \\sim E_K \\sim t^{-1.1}\\).  Length scales in the magnetic field increase as \\(t^{0.5}\\) in the helical case, but slightly slower in the non-helical case, \\(t^{0.4}\\). Note that these growth laws disagree with all theoretical predictions to date \\cite{Bis93,Ole97,Field:2000hi,Son:1999my}, which give $t^{2/3}$ in the helical case and $t^{2/7}$ for our power spectrum in the non-helical case.  Helical magnetic fields are found to evolve in a self-similar way, with a scaling function \\(g_M(z) \\sim z^{-p}\\) at large \\(k\\), where \\(p=2.5\\) for \\(\\mathrm{Re}\\sim10^2\\). Note that this is significantly different from both the Iroshnikov-Kraichnan and Kolmogorov spectra, $k^{-3/2}$ and $k^{-5/3}$, respectively.  A good theoretical understanding of these scaling laws is required before the evolution of magnetic fields in the early Universe is properly understood, as a small error in the exponent makes a large error in the prediction of the magnetic field strength when propagated over many orders of magnitude in time.  For example, Vachaspati's contribution to these proceedings \\cite{Vac01} assumes a growth law of $t^{2/3}$ in the length scale, based on a simple argument invoking helicity conservation \\cite{Bis93,Son:1999my,Field:2000hi}, to obtain seed fields of an interesting strength from the electroweak transition, which is quite different from our observed growth law of $t^{0.5}$."
    },
    "0201/astro-ph0201520_arXiv.txt": {
        "abstract": "We study the angular momentum profile of dark matter halos for a statistical sample drawn from a set of high-resolution cosmological simulations of $256^3$ particles.  Two typical Cold Dark Matter (CDM) models have been analyzed, and the halos are selected to have at least $3\\times 10^4$ particles in order to reliably measure the angular momentum profile. In contrast with the recent claims of Bullock et al., we find that the degree of misalignment of angular momentum within a halo is very high. About 50 percent of halos have more than 10 percent of halo mass in the mass of negative angular momentum $j$. After the mass of negative $j$ is excluded, the cumulative mass function $M(<j)$ follows approximately the universal function proposed by Bullock et al., though we still find a significant fraction of halos ($\\sim 50\\%$) which exhibit systematic deviations from the universal function. Our results, however, are broadly in good agreement with a recent work of van den Bosch et al..  We also study the angular momentum profile of halos in a Warm Dark Matter (WDM) model and a Self-Interacting Dark Matter (SIDM) model. We find that the angular momentum profile of halos in the WDM is statistically indistinguishable from that in the CDM model, but the angular momentum of halos in the SIDM is reduced by the self-interaction of dark matter. ",
        "introduction": "The formation of galactic disks is one of the most significant unsolved problems in cosmology. In the popular hierarchical clustering framework (White \\& Reese 1978), disks form in the potential wells of dark matter halos while the baryons cool and collapse dissipatively (Fall \\& Efstathiou 1980). Assuming that the angular momentum of the gas is conserved during the collapse and that the gas has the same spin $\\lambda$ parameter as the dark matter in a halo, one can derive a density profile of cooled gas under an additional assumption that either the density profile is assumed to be exponential as observed or the angular momentum of dark matter within a halo follows a distribution, say, given by a uniform sphere in solid body rotation (e.g. Fall \\& Efstathiou 1980, Blumenthal et al.1986, Dalcanton, Spergel, \\& Summers 1997, Jimenez et al.1997, Mo, Mao, \\& White 1998; van den Bosch 1998; Avila-Reese \\& Firmani 2000). While the theory is quite successful in explaining many observational data of spiral galaxies (e.g. Mo, Mao, \\& White 1998), it has encountered two potentially serious difficulties.  It has been known for many years that detailed numerical simulations of gas collapse in Cold Dark Matter (CDM) models (Navarro \\& Benz 1991; Navarro \\& White 1994; Steinmetz \\& Navarro 1999; Navarro \\& Steinmetz 2000) have consistently indicated that the infalling gas loses too much angular momentum due to the dynamical friction, and the resulting disks are accordingly too small to be compatible with the observations.  This discrepancy is known as the ``angular momentum catastrophe'' in disk galaxy formation(Navarro, \\& Benz 1991). It is unknown if this is a failure of CDM models (Sommer-Larsen \\& Dolgov 2001) or this is due to the simplified treatments of complicated gas physics in the simulations, e.g. gas heating from supernova explosion (Weil et al. 1998; Sommer-Larsen, Gelato \\& Vedel 1999; Mo \\& Mao 2002).  The second difficulty emerged recently after Bullock et al. (2001; hereafter B2001) claimed, based on a high-resolution simulation of a low density flat CDM model (LCDM), that the dark matter halos have a universal angular momentum profile, i.e. the mass distribution of specific angular momentum $j$ in a halo obeys a universal form, \\begin{equation} M(<j)= \\frac{M_{\\rm v}\\mu j}{j_{\\rm 0}+ j} \\, \\label{eq:jdis} \\end{equation} where $M_{\\rm v}$ is the virial mass, $ j_{\\rm 0}=(\\mu-1) j_{\\rm max}$, and $j_{\\rm max}$ is the maximum specific angular momentum. The parameter $\\mu$ indicates the shape of the profile.  Under the conventional assumption that the angular momentum of the gas is the same as that of the dark matter before gas collapse and is conserved in detail during the collapse, the distribution (\\ref{eq:jdis}) implies that the formed disks contain too much low angular momentum mass to be compatible  with the high-resolution observations of galactic disks(Bullock et al. 2001, van den Bosch 2001; van den Bosch et al. 2001).  In this work we will examine the validity of equation (\\ref{eq:jdis}) using independent simulation data. A large set of cosmological N-body simulations of the Standard Cold Dark Matter (SCDM) model and a LCDM model (Jing \\& Suto 1998) are employed. The simulations have a similar resolution to that of B2001, but the sample used in this work is significantly larger. The large sample enables us to study the angular momentum profile by selecting a sample of halos which are well resolved. The two cosmological models will help to answer if the angular momentum profile of halos depends on cosmological parameters. We will also analyze the individual halo simulations (Jing \\& Suto 2000) in a Warm Dark Matter (WDM) model (Schaeffer \\& Silk 1988, Colombi, Dodelson \\& Widrow 1996, Sommer-Larsen \\& Dolgov 2001, Jing 2001, Avila-Reese et al.2001, Bode, Ostriker, \\& Turok 2001) and in a Self-Interaction Dark Matter (SIDM) model (Spergel \\& Steinhardt 2000, Burkert 2000, Dav{\\'e} et al. 2001, Yoshida et al. 2000) to study if the angular momentum profile of halos is different in different dark matter models. These alternative dark matter models have been recently proposed to solve possible problems facing CDM models at sub-galactic scales, i.e. central steep density profiles of halos and too numerous sub-halos (Moore et al. 1999, Klypin et al. 1999).  Several similar works had appeared when the present work was near completion. Knebe et al. (2001) and Bullock et al.(2001b) have studied the angular momentum profile of halos in warm dark matter models which are similar to (but not the same as) ours. van den Bosch et al. (2002) also examined whether equation (\\ref{eq:jdis}) is a good prescription for the angular momentum profile of halos in an N-body/hydrodynamic simulation. Since the simulations used by these authors are quite distinct, in terms of simulation volumes and simulation techniques, from those used in this paper, our work interestingly complements to theirs. Wherever necessary, we will compare with their results in the following sections.  The outline of this paper is as follows. The cosmological models and the numerical simulations are described in the next section.  The methods for measuring the angular momentum profile and the results for the cosmological simulations are presented in section 3. In section 4, we discuss the results for the WDM and SIDM models. The conclusions are given in section 5. ",
        "conclusions": "We have studied the angular momentum profile of dark matter halos for a statistical sample drawn from a set of high-resolution cosmological simulations.  Two typical CDM models have been considered, and the halos are selected to have at least $3\\times10^4$ particles in order to reliably measure the angular momentum profile. In contrast with the recent claims of B2001, we find that the degree of misalignment of the angular momentum within a halo is very high. About 50 percent of halos have more than 10 percent of halo mass in the mass of negative angular momentum $j$. After the mass of negative $j$ is excluded, the cumulative mass function $M(<j)$ follows approximately the universal function proposed by B2001 if the virial mass $M_{\\rm v}$ in the function is replaced by $M(j>0)$, though we still find a significant fraction of halos ($\\sim 50\\%$) which exhibit systematic deviations from the universal function. Our results, however, are broadly in good agreement with the recent work of van den Bosch et al. (2002).  We have also studied the angular momentum profile of halos in the Warm Dark Matter model and in the Self-Interacting Dark Matter model in order to study how the angular momentum profile is affected by the basic assumption about the dark matter. We have made a detailed comparison between the halos in these scenarios and the corresponding halos in the LCDM model. We find that there is no {\\it systematic} difference in the angular momentum between the halos from the WDM and from the LCDM, though a pair of corresponding WDM and LCDM halos may exhibit quite different angular momentum profiles. We also find that the self-interaction of dark matter in the SIDM models can generally reduce the angular momentum, which makes the spin parameter $\\lambda$ and the shape parameter $\\mu$ smaller. Thus it seems that these dark models do little help to solve the angular momentum problem encountered by the CDM models.  Our results also indicate that it should be cautious to use the universal angular momentum profile of B2001 to predict observational properties for disk galaxies. The angular momentum in different parts of a halo does not orient as coherently as B2001 claimed. The mass of negative angular momentum $j$ may combine with those mass of small $j$ to form the bulge component in spiral galaxies (van den Bosch et al. 2002), or the angular momentum profile of the gas in a halo is significantly different from that of the dark matter due to hydro-dynamical processes like heating and explosions (Maller \\& Dekel 2002). The relation between the disk properties and the angular momentum profile (Eq.\\ref{eq:jdis}) could be more complicated than previously thought."
    },
    "0201/astro-ph0201534_arXiv.txt": {
        "abstract": "We present the first non-local ($z>$0.2) measurement of the cluster-cluster spatial correlation length, using data from the Las Campanas Distant Cluster Survey (LCDCS).  We measure the angular correlation function for velocity-dispersion limited subsamples of the catalog at estimated redshifts of $0.35\\le z_{est}<0.575$, and derive spatial correlation lengths for these clusters via the cosmological Limber equation.  The correlation lengths that we measure for clusters in the LCDCS are consistent both with local results for the APM cluster catalog and with theoretical expectations based upon the Virgo Consortium Hubble Volume simulations and the analytic predictions.  Despite samples containing over 100 clusters, our ability to discriminate between cosmological models is limited because of statistical uncertainty. ",
        "introduction": "The spatial correlation function of galaxy clusters provides an important cosmological test, as both the amplitude of the correlation function and its dependence upon mean intercluster separation are determined by the underlying cosmological model. In hierarchical models of structure formation, the spatial correlation length, $r_0$, is predicted to be an increasing function of cluster mass, with the exact dependence determined by $\\sigma_8$ (or equivalently $\\Omega_0$, using the constraint on $\\sigma_8-\\Omega_0$ from the local cluster mass function) and the power spectrum shape parameter, $\\Gamma$.  Low density and low $\\Gamma$ models generally predict stronger clustering for a given mass and a greater dependence of the correlation length upon cluster mass.  The three-space correlation function of clusters was first measured for subsamples of the Abell catalog by \\citet{bah83} and \\citet{kly83}. Both groups found that the correlation function is well-described by a power law, $\\xi(r)=(r/r_0)^{-\\gamma}$, and obtained a correlation length $r_0\\simeq 25 h^{-1}$ Mpc with $\\gamma$$\\simeq$2.  \\citet{bah83} also observed a strong dependence of correlation strength upon cluster richness, which was later quantified by \\citet{bah88} and \\citet{bah92w} as a roughly linear dependence of $r_0$ upon $d_c$, the mean intercluster separation. \\citet{pos92} and \\citet{pea92} confirmed the form of the correlation function for the Abell catalog in their larger spectroscopic samples, with both studies obtaining $r_0$$\\simeq$20$h^{-1}$ Mpc for clusters with richness class R$\\ge$1.  While these correlation lengths have strong implications for cosmological models, a key problem with interpretation of the Abell results is concern that the observed correlation lengths are positively skewed by projection effects and sample inhomogeneities \\citep[see][]{sut88,dek89,efs92,pea92}. Several analyses find that $\\xi(\\sigma,\\pi)$ is strongly anisotropic, evidence that these effects are significant \\citep{sut88,efs92,pea92}. Still, the net impact of these factors is unclear. Contrary to the concerns raised by these studies, \\citet{mil99} use an expanded sample of Abell clusters to derive correlation lengths that are consistent with earlier analyses and robust to projection effects, and \\citet{van97} argue that projections are insufficient to account for the stronger correlation observed in the Abell catalog as compared to the APM catalog \\citep{dal92}.   Fortunately, independent constraints on the correlation function have arisen as new catalogs with automated, uniform selection criteria have become available \\citep[e.g.,][]{dal92,nic92,dal94,nic94,rom94,cro97,collins2000,mos2000}. Some of the recent optical catalogs, such as the APM \\citep{dal92}, also probe to lower $d_c$ than the Abell samples, while the X-ray samples provide greatly improved leverage for the most massive (highest $d_c$) clusters. While systematic variations persist between samples, the existing data are generally consistent with $r_0$ slowly increasing with $d_c$ \\citep[however, see][]{collins2000}.  To exploit the growth of observational data, there has been a corresponding theoretical effort to predict the cluster spatial correlation function as a function of mass and epoch. Driven by cosmological volume N-body simulations \\citep[e.g.,][]{gov99,col2000,mos2000} and the development of a well-tested analytic formalism \\citep{mo96,she99,smt1999}, a theoretical framework has been established that enables derivation of quantitative cosmological constraints from the observational data. The mass dependence can be studied using existing data sets and is typically best matched by low-density models \\citep[see][]{cro97,bor99b,collins2000}, although systematic uncertainties and observational scatter have precluded precision cosmological constraints. In contrast, the redshift dependence remains unconstrained because the data have not existed to test the evolutionary predictions of these models.  In this paper we utilize the Las Campanas Distant Cluster Survey (LCDCS) to determine the spatial correlation length at $z$$\\simeq$0.45. We first measure the angular correlation function for a series of subsamples at this epoch and then derive the corresponding $r_0$ values via the cosmological Limber inversion \\citep{peebles80,efs91,hud96}. The resulting $r_0$ values constitute the first measurement at this epoch of the dependence of the cluster correlation length upon $d_c$, probing mean separations similar to previous local optical catalogs. Popular structure formation models predict only a small amount of evolution from z=0.45 to the present, as illustrated in section 5. We test this prediction by comparing our results with the local observations. ",
        "conclusions": "\\label{sec-cosmodiscussion}  The Las Campanas Distant Cluster Survey is the largest existing catalog of clusters at $z$$>$0.3, providing a unique sample with which to study the properties of the cluster population.  We have used the LCDCS to constrain the cluster-cluster angular correlation function, providing the first measurements for a sample with a mean redshift $z$$\\ga$0.2.  From the observed angular correlation function, we derive the spatial correlation length, $r_0$, as a function of mean separation, $d_c$. We find that the LCDCS correlation lengths are in agreement with results from local samples, and observe a dependence of $r_0$ upon $d_c$ that is comparable to the results of \\citet{cro97} for the APM catalog.  This clustering strength, its dependence on number density, and its minimal redshift evolution are consistent with analytic expectations for low density models, and with results from the $\\Lambda$CDM Hubble Volume simulations. Consequently, while statistical uncertainty limits our ability to discriminate between cosmological models, our results are in concordance with the flat $\\Lambda$CDM model favored by recent supernovae and cosmic microwave background observations \\citep[e.g.,][]{riess01,pryke01,boom01}."
    },
    "0201/astro-ph0201228_arXiv.txt": {
        "abstract": "We present photometry of 12 recent supernovae (SNe) recovered in a {\\it Hubble Space Telescope} Snapshot program, and tie the measurements to earlier ground-based observations, in order to study the late-time evolution of the SNe. Many of the ground-based measurements are previously unpublished, and were made primarily with a robotic telescope, the Katzman Automatic Imaging Telescope.  Evidence for circumstellar interaction is common among the core-collapse SNe.  Late-time decline rates for Type IIn SNe are found to span a wide range, perhaps due to differences in circumstellar interaction. An extreme case, SN IIn 1995N, declined by only 1.2 mag in $V$ over about 4 years following discovery. Template images of some SNe must therefore be obtained many years after the explosion, if contamination from the SN itself is to be minimized. Evidence is found against a previous hypothesis that the Type IIn SN 1997bs was actually a superoutburst of a luminous blue variable star. The peculiar SN Ic 1997ef, a ``hypernova,\" declined very slowly at late times. The decline rate of the SN Ia 2000cx decreased at late times, but this is unlikely to have been caused by a light echo. ",
        "introduction": "Supernovae (SNe) represent the final, explosive stage in the evolution of certain varieties of stars (see, e.g., Woosley \\& Weaver 1986; Arnett \\etal~1989; Wheeler \\& Harkness 1990; Filippenko 1997, for reviews of SN types and explosion mechanisms). They synthesize and expel heavy elements, heat the interstellar medium, trigger vigorous bursts of star formation, create neutron stars and sometimes black holes, and produce energetic cosmic rays. Type Ia SNe (SNe Ia), among the most luminous of SNe, are exceedingly useful cosmological tools and have been used to study the expansion history of the Universe.  These studies (Riess \\etal~1998, 2001; Perlmutter \\etal~1999; see Filippenko 2001 for a recent summary) reveal the surprising result that the expansion of the Universe is currently accelerating, perhaps due to a nonzero cosmological constant. SNe are clearly among the most interesting and important constituents of the Universe and should be vigorously studied.  \\subsection{Discovery and Monitoring of Nearby Supernovae}  Detailed spectral and photometric observations of SNe can be used to study the properties of SNe, compare SNe at different redshifts, and gain insights into the evolutionary paths that lead to these energetic explosions.  However, until recently, most relatively nearby SNe were found either sporadically, in images taken for other purposes, or at a considerable time after the explosion. Moreover, systematic follow-up observations were scarce for all but a minor fraction of the SNe. During the past few years, however, the situation has changed dramatically, as robotic (or nearly robotic) telescopes have been dedicated to the search for and follow-up of SNe (Filippenko \\etal~2001).  The two outstanding examples are the Beijing Astronomical Observatory Supernova Search (BAOSS; Li \\etal~1996) with a 0.6-m telescope, and the Lick Observatory Supernova Search (LOSS; Treffers \\etal~1997; Li \\etal~2000; Filippenko \\etal~2001) with the 0.8-m Katzman Automatic Imaging Telescope (KAIT).  KAIT reaches a limit of $\\sim$ 19 mag (4$\\sigma$) in 25-s unfiltered, unguided exposures, while 5-min guided exposures yield $R \\approx $ 20 mag. 2500$-$7000 galaxies (most with $cz \\leq$ 6000 km s$^{-1}$) are surveyed during any particular season, with a cycle time of $\\sim$ 3$-$10 days. The search software automatically subtracts new images from old ones and identifies SN candidates which are subsequently examined by undergraduate research assistants.  LOSS discovered its first SN in 1997 (SN 1997bs; Treffers \\etal~1997; Van Dyk \\etal~2000), then 20 SNe in 1998, 40 in 1999, 36 in 2000, and 68 in 2001. Together, LOSS and BAOSS have found a total of over 180 SNe in the past 5 years. Recently, LOSS teamed up with Michael Schwartz in Arizona, forming the Lick Observatory and Tenagra Observatory Supernova Search (LOTOSS), allowing more SNe to be discovered and followed.  Since most of the SNe found during these systematic searches are discovered very early in their development, and have a large amount of follow-up time devoted to them, they are among the world's best-studied SNe (see, e.g., Li \\etal~2000, 2001a; Leonard \\etal~2002a,b). Li, Filippenko, \\& Riess (2001) and Li \\etal~(2001b) also show that these surveys are least affected by observational biases, and provide the most accurate luminosity function for SNe. The rate of intrinsically peculiar SNe~Ia, for example, is found to be unexpectedly high.  \\subsection{The Environment and Late-Time Detection of Supernovae}  The explosion of a SN leaves few traces of the star that underwent the catastrophic event. An important clue to the nature of the progenitor is its environment (e.g., Boffi 1999). Unfortunately, most studies of the sites of SNe have been hampered by the limited spatial resolution of ground-based observations (e.g., van den Bergh 1988; Panagia \\& Laidler 1991; Boffi, Sparks, \\& Macchetto 1999). This problem can be partially overcome by using {\\it Hubble Space Telescope (HST)} archival images (e.g., Van Dyk \\etal~1999a,b, 2000). However, although the chance that the site of any particular SN has been imaged by {\\it HST} during other programs is rapidly growing, it is still relatively low. Also, some of the positions of older SNe (prior to 1990) were uncertain, compromising studies of their local environment.  To remedy this, an {\\it HST} Snapshot survey program (GO-8602) of recent, nearby SNe was conducted starting in Cycle 9, requesting WFPC2 images of the sites of 23 LOSS and BAOSS SNe\\footnote{Some of these 23 SNe were first discovered by other groups, but subsequently found in the course of LOSS and BAOSS.} with $cz \\leq $6000 km s$^{-1}$.  These SNe provide the best spatial resolution and have accurate positions (often better than $\\pm 0\\farcs 5$). In addition, they were generally discovered early in their development, and had extensive ground-based follow-up studies. 45 out of the 70 ($\\sim 45\\%$) requested observations have been made, with 20 out of the 23 SNe ($\\sim 90\\%$) observed at least once and 13 ($\\sim 56\\%$) of the SNe having images in more than one filter; thus, the chances of imaging the sites of specific SNe are much higher than by using the archive alone. Moreover, the data are all obtained in a uniform way using the same set of filters.  One advantage of concentrating the Snapshot survey on very recent, nearby, well-studied SNe is that many of the SNe themselves are still visible in the $HST$ images, providing late-time photometry superior to that achieved from the ground; at such late times, the SNe are so faint that their ground-based photometry is contaminated by neighboring stars within the seeing disk. Late-time photometry, especially through more than one filter, provides not only useful information on the underlying physics for the lingering light, such as radioactive decay of long-lived isotopes, interaction with circumstellar matter, and light echoes, but also the relative brightness of SNe still present in ``template\" images taken at various epochs compared to their maximum light. To obtain proper photometry of a SN, which often occurs in a complicated background (e.g., spiral arms or H~II regions), observers are required to take template images of the host galaxy a year or two after the discovery, and then do image subtraction. It is often assumed that the light of the SN is essentially negligible in these template images, but as we later discuss in this paper, some SNe (especially those of Type II) are quite long-lived and may contaminate the template images.  Detailed analysis of the SN environments is still underway, and the results will be discussed elsewhere.  In this paper, we report on SNe recovered in the Snapshot images. Section 2 contains a description of the observations and analysis of the photometry, while \\S~3 presents the light curves of all recovered SNe, details of individual SNe, and comparison of late-time light-curve shapes and decline rates. Our conclusions are summarized in \\S~4. ",
        "conclusions": "In this paper we present 12 SNe recovered in the $HST$ Cycle 9 Snapshot program GO-8602. $HST$ photometry was obtained for the SNe, and the data were tied to early-time ground-based observations, in order to study the late-time evolution of various types of SNe. Much of the ground-based data presented here are previously unpublished, and were obtained primarily with the Katzman Automatic Imaging Telescope.  The successful detection of SN 1996cb, a SN~IIb, more than 4 years after discovery suggests the existence of interaction between the SN ejecta and the circumstellar medium.  The peculiar SN~Ic 1997ef, a ``hypernova,\" is found to decline slowly between 100 and 1200 days after discovery.  Two bright SN II-P discovered in 1999, SNe 1999gi and 1999gq, are also both recovered in the Snapshot observations.  We find that there is a diversity in the late-time evolution of SNe~IIn, perhaps due to differences in circumstellar interaction. Some decline monotonically, such as SN 1999eb within the first 660 days following maximum brightness. Others experience multiple changes in their decline rate, such as SNe 1998S, 1999el, and 2000P. SN 1995N fades very slowly, declining by only 1.2 mag in $V$ over $\\sim 1400$ days following discovery. Thus, template images of some SNe must be obtained many years after the explosion, to minimize contamination from the SN itself.  SN 1997bs, a subluminous SN~IIn that was suggested by Van Dyk \\etal (2000) to be a superoutburst of a massive luminous blue variable, is only marginally detected in the F555W observations and not at all in the two F814W images, casting some doubt on the hypothesis that the progenitor survived the explosion.  SN 1999bw, an object similar to SN 1997bs in terms of its spectrum and peak luminosity, showed a slower late-time decline rate than SN 1997bs.  SN 2000cx is the only SN~Ia recovered in the Snapshot observations.  Both the $R$-band and $I$-band decline rates are found to decrease at late times, but we do not believe this is due to a light echo.  Circumstellar interaction is apparently very common in the late-time evolution of core-collapse SNe, such as SNe~IIn, SNe~IIb, and possibly some SNe~II-P and SNe~Ic."
    },
    "0201/astro-ph0201158_arXiv.txt": {
        "abstract": "We study the pulse morphologies and pulse amplitudes of thermally emitting neutron stars with ultrastrong magnetic fields.  The beaming of the radiation emerging from a magnetar was recently shown to be predominantly non-radial, with a small pencil and a broad fan component. We show that the combination of this radiation pattern with the effects of strong lensing in the gravitational field of the neutron star yields pulse profiles that show a qualitatively different behavior compared to that of the radially-peaked beaming patterns explored previously.  Specifically, we find that: {\\em (i)} the pulse profiles of magnetars with a single hot emission region on their surface exhibit $1-2$ peaks, whereas those with an antipodal emission geometry have $1-4$ peaks, depending on the neutron star compactness, the observer's viewing angle, and the size of the hot regions; {\\em (ii)} the energy dependence of the beaming pattern may give rise to weakly or strongly energy-dependent pulse profiles and may introduce phase lags between different energy bands; {\\em (iii)} the non-radial beaming pattern can give rise to high pulsed fractions even for very relativistic neutron stars; {\\em (iv)} the pulsed fraction may not vary monotonically with neutron star compactness; {\\em (v)} the pulsed fraction does not decrease monotonically with the size of the emitting region; {\\em (vi)} the pulsed fraction from a neutron star with a single hot pole has, in general, a very weak energy dependence, in contrast to the case of an antipodal geometry.  Comparison of these results to the observed properties of anomalous X-ray pulsars strongly suggests that they are neutron stars with a single hot region of ultrastrong magnetic field. ",
        "introduction": "Anomalous X-ray Pulsars (AXPs) and Soft Gamma-ray Repeaters (SGRs) are two classes of intriguing objects that have challenged the standard paradigm of young neutron stars. Particularly in the case of AXPs, two types of models have been proposed to explain their spectral and timing properties: conventional accretion models modified to account for the absence of an observable donor star (van Paradijs, Taam, \\& van den Heuvel 1995; Chatterjee, Hernquist, \\& Narayan 2000; Alpar 2001) and magnetar models with several different mechanisms for powering the X-ray emission (Thompson \\& Duncan 1996; Heyl \\& Hernquist 1998). In all these magnetar models, a significant fraction of the X-ray emission is thought to originate at the neutron star surface.  Recently, there has been significant progress in the calculation of the surface emission properties of ultramagnetized neutron stars and their application to AXPs (\\\"Ozel 2001; \\\"Ozel, Psaltis, \\& Kaspi 2001; see also Ho \\& Lai 2001; Zane et al.\\ 2001). \\\"Ozel (2001) showed that the angle dependence of surface radiation in such strong fields has both a pencil and a fan component, with the latter dominating in the X-ray range. Such a beaming pattern can give rise to a variety of pulse profiles and timing properties that have not been explored to date. Furthermore, the general relativistic bending of photon trajectories in the strong gravitational field of the neutron star significantly affects the observed properties of emission originating at the stellar surface. This is in contrast to rotation-powered pulsars whose radio emission originates at the light cylinder far above the stellar surface, where the self lensing by the neutron star is weak. Therefore, in order to compare in detail the models with the current and future observations of AXPs, it is now necessary to carry out a complete study of the observable properties of magnetars taking all the above effects into account.  The advent of X-ray telescopes with good timing and energy resolution as well as broad spectral coverage has produced high-quality data on AXPs in the recent years, making such a study timely. Indeed, observations with {\\em ASCA} and {\\em BeppoSAX} have yielded good broad-band spectra of AXPs as well as a measure of the flux variations during a pulse cycle in multiple energy bands (e.g., White et al.\\ 1996; Oosterbroek et al.\\ 1998). Moreover, the superb timing resolution of the {\\em Rossi X-ray Timing Explorer (RXTE)} made possible the phase connection of pulse cycles over several years and hence produced detailed energy-dependent pulse profiles (Gavriil \\& Kaspi 2002). Finally, the grating spectrographs onboard the {\\em Chandra} and {\\em XMM-Newton} observatories extend the energy range of current observations towards the soft X-rays in addition to allowing for phase-resolved spectroscopy and searches for spectral lines (e.g., Patel et al.\\ 2001; Juett et al.\\ 2002; Tiengo et al.\\ 2002). It has been argued that the combination of the pulse profiles, the amplitude of pulsations, and the pulse-phase resolved spectral features observed with these instruments provide the most stringent constraints on the models of AXPs (\\\"Ozel et al. 2001).  In this paper we explore the pulse profiles expected from a strongly magnetized neutron star emitting thermally from its hot surface.  In detail, we focus on the dependence of the pulse profiles on the emission geometry, the orientation of the observer, the photon energy, and the compactness of the neutron star. We also discuss the implications of our results for magnetar models of AXPs. ",
        "conclusions": "We studied the timing properties of ultramagnetized neutron stars emitting thermally from their surfaces. We used the energy and angle dependence of the emerging radiation obtained from the recent detailed calculations of radiative equilibrium atmospheres in ultrastrong magnetic fields (\\\"Ozel 2001, 2002). We combined them with the general relativistic calculation of photon transport from the neutron star surface to an observer at infinity and calculated the expected pulse profiles for a wide range of model parameters. We found that the combination of the non-radially peaked beams relevant to magnetars with the strong gravitational lensing leads to a number of qualitatively new and interesting results on the pulse morphology and amplitudes, which we summarize below.  \\noindent {\\bf 1.} An emission geometry consisting of one hot pole gives rise to one or two peaks per pulse cycle, whereas an antipodal geometry can produce one to four peaks, depending strongly on the orientation angles $\\alpha$ and $\\beta$ as well as on the relativity parameter $p$ and the angular size of the emitting region $\\rho$.  \\noindent {\\bf 2.} The pulse profiles in different energy bands show a wide range of morphologies in both emission geometries.  \\noindent {\\bf 3.} The non-radial beaming of the thermal radiation emerging from an ultramagnetized neutron-star surface can give rise to high pulsed fractions even for very compact neutron stars.  \\noindent {\\bf 4.} In the case of an antipodal emission geometry, the pulsed fraction may not be a monotonic function of the relativity parameter $p$, in contrast to the case of radially peaked beaming patterns.  \\noindent {\\bf 5.} The pulsed fraction does not decrease monotonically with the size of the emitting region but shows a secondary maximum at intermediate ($\\sim 60-70^\\circ$) angular sizes.  \\noindent {\\bf 6.} The pulsed fraction in the antipodal emission geometry shows a characteristic increase with photon energy in the 1--8~keV range. In contrast, the pulsed fraction in the single pole case has, in general, a very weak energy dependence.  Note that a number of simplifying assumptions have been made in the models presented here. We have taken the temperature across a hot region to be uniform, which may not be realistic for all the mechanisms that could power the surface emission. In the case of a cooling neutron star with a dipole magnetic field, the temperature is a function of the magnetic latitude and thus shows variations across the surface of the neutron star (e.g., Page 1995). Also in the case of a single hot pole, similar variations can also be expected. This can affect the pulse profiles, possibly reducing the pulsed fractions and smoothing out the pulse morphologies. Similarly, allowing for some cooler emission from the rest of the neutron star can alter the pulse profiles and may account for the additional complexity observed in the energy-dependent pulse profiles of AXPs.  We also assumed that the emission properties are determined entirely by the stellar atmosphere and have not taken into account the possible effects of the neutron star's magnetosphere further above (see, e.g., Thompson, Lyutikov, \\& Kulkarni 2002).  At present, the structure of the magnetosphere of a strongly magnetized neutron star is not well understood and can only be treated through parametrizations. Note that the alternative class of models of AXPs that rely on accretion onto a magnetized neutron star also require the study of the processes that take place in the accretion column in the neutron star's magnetic field. It was shown earlier that such processes can also lead to a variety of pulse profiles (e.g., M\\'esz\\'aros \\& Nagel 1985) but further study is necessary to determine their relevance for AXPs.  The results presented above have direct implications for the thermally emitting magnetar models of AXPs. To carry out a comparison with observations, we summarize the timing properties of these sources which can be used to constrain such models.  First, the observed pulse profiles show two prominent peaks per pulse cycle in four sources and a single peak in the fifth source (e.g., Gavriil \\& Kaspi 2002). Second, the energy dependence of the pulse morphology can be very weak, as in 1E~1048.1$-$5937, or quite strong, as in 1RXS~1708$-$4009 (Gavriil \\& Kaspi 2002). Third, AXPs can have pulsed fractions as high as 70\\%, together with luminosities that are high for their low inferred effective temperatures (\\\"Ozel et al.\\ 2001). Finally, their pulsed fractions show a weak dependence on photon energy (e.g., Oosterbroek et al.\\ 1998).  All the above properties are hard to account for in a thermally cooling magnetar model (e.g., Heyl \\& Hernquist 1998) that has a two-fold symmetry and thus has pulse profiles very similar to those of the antipodal emission geometry. In particular, the number of observed peaks in the pulse profiles as well as the strong energy-dependence predicted in the two-pole models do not reproduce the observations. Furthermore, the pulsed fractions in such a model are in general lower than the values observed in AXPs.  The observations, on the other hand, strongly suggest that AXPs are neutron stars with a single, hot region of ultrastrong magnetic field. A single hot pole can be realized in a number of ways. It may be powered by processes such as crustal cracking, magnetic-field reconfiguration, or decay of magnetic multipoles (e.g., Thompson \\& Duncan 1996). It could also arise from an off-centered magnetic dipole that renders one of the poles either not observable or else very close to the other pole.  All of these possibilities need further theoretical investigations in order to determine the geometry of emission, the total energy output and the lifetime of the temperature asymmetries on the surface.  In addition, because the first set of processes are expected to occur randomly, they can produce time-variable pulse profiles. Observations of pulse morphology changes similar to the {\\em GINGA} observation of 1E~2259$+$586 (Iwasawa, Koyama, \\& Halpern 1992) can help constrain the mechanism powering the thermal emission of AXPs. Finally, investigating the energy dependence of the pulsed fractions in the soft X-rays ($\\lesssim 1$~keV) with {\\em Chandra\\/} and {\\em XMM-Newton\\/} as well as in longer wavelengths will extend the baseline over which models can be compared to data and provide the most stringent constraints."
    },
    "0201/astro-ph0201191_arXiv.txt": {
        "abstract": "High speed photometry of KUV 01584-0939 (alias Cet3) shows that it has a period of 620.26~s. Combined with its hydrogen-deficient spectrum, this implies that it is an AM CVn star. The optical modulation is probably a superhump, in which case the orbital period will be slightly shorter than what we have observed. ",
        "introduction": "The star KUV 01584-0939 (hereafter referred to as Cet3, as in the Downes, Webbink \\& Shara (1997) catalogue of Cataclysmic Variables (CVs)) was discovered during the Kiso survey (Kondo, Noguchi \\& Maehara 1984) for ultraviolet rich objects. A spectrum of Cet3 covering the wavelength range 4000 -- 7100 {\\AA} was obtained by Wegner, McMahan \\& Boley (1987) who drew attention to the great strength of He\\,II emission and the weakness of the Balmer emission lines. We have observed Cet3 as part of a high time resolution photometric survey of faint CVs (Woudt \\& Warner 2001). ",
        "conclusions": "AM CVn stars evolve from normal composition binaries which have experienced two phases of mass exchange, exposing the helium cores of the original stars. During their orbital evolution such systems pass through a minimum period near 4 mins and thereafter evolve to longer periods (Tutukov \\& Yungelson 1996). The driving mechanism for orbital evolution in such short period binaries is loss of angular momentum through emission of gravitational radiation (Paczynski 1967). This determines the rate of mass transfer, $\\dot{M}$, and predicts (Warner 1995b) values $\\sim 1 \\times 10^{-8}$ M$_\\odot$ y$^{-1}$ at $P_{orb}$ = 10 min, decreasing steeply to $4 \\times 10^{-11}$ M$_\\odot$ y$^{-1}$ at $P_{orb}$ = 40 min. These rates are compatible with the range of observed properties of the seven optically observed (Warner 1995b; Ruiz et al.~2001) AM CVn stars and the X-ray system (Cropper et al.~1998). The latter (V407 Vul = RX\\,J1914+24) has $P_{orb}$ = 569 s, is thought to be a polar (i.e., the primary is strongly magnetic, so its rotation is synchronised with the revolution of the secondary), but is very faint optically because it is obscured by an interstellar cloud. In contrast, Cet3 has a similar period and is relatively accessible at $m_V = 16.9$.  If Cet3 has $\\dot{M} \\sim 1 \\times 10^{-8}$ M$_\\odot$ y$^{-1}$, its accretion luminosity will be $\\sim$ 5 L$_\\odot$ for a 0.7 M$_\\odot$ white dwarf primary. Most of this will be radiated at the inner boundary layer of the accretion disc at a temperature of $\\sim 3 \\times 10^5$ K, so Cet3 should be a soft X-ray source of high intrinsic luminosity. This also accounts for the strength of the helium and carbon ionic spectra in Cet3. We calculate that d{$P_{orb}$}/d$t$ $\\sim 6 \\times 10^{-12}$, which implies that d{$P_{orb}$}/d$t$ should be measurable in only a few years of observations if the modulation indeed has an orbital origin. We have no means at present to estimate the distance of Cet3, but we expect it to be a strong source of gravitational radiation (Warner 1995b; Hils \\& Bender 2000).  We can estimate parameters for Cet3 in the following way. As the secondary must fill its Roche lobe, a $P_{orb}$ of 620 s implies a mean density $\\bar{\\rho} \\simeq 3.6 \\times 10^3$ gm cm$^{-3}$ (see equation 2.3b of Warner 1995a). From the mass-radius relation for low mass white dwarfs (Tutukov \\& Yungelson 1996), this gives a mass for the secondary $M(2)$ = 0.066 M$_\\odot$. For an assumed primary mass $M(1)$ = 0.7 M$_\\odot$, we have $q = M(2)/M(1)$ = 0.094 and a binary separation of $9.9 \\times 10^9$ cm. Cet 3 is a very compact system: from standard formulae (Warner 1995a) we find an accretion disc radius $\\sim 5.4 \\times 10^9$ cm, the radius $R(2)$ of the secondary $\\sim 2.1 \\times 10^9$ cm, and the radius of the primary $R(1) \\sim 0.85 \\times 10^9$ cm. The accretion disc therefore extends only $\\sim 5.4 R(1)$ above the surface of the primary and will be strongly irradiated by the $\\sim 3 \\times 10^5$ K source at the inner edge of the disc.  In addition, for $\\dot{M} \\sim 1 \\times 10^{-8}$ M$_\\odot$ y$^{-1}$ the maximum temperature in the disc itself is $\\sim 6.2 \\times 10^4$ K. The surface of the secondary intercepts $\\sim$ 1\\% of the soft X-ray emission and, with allowance for shielding by the accretion disc, the side facing the primary will as a result be heated to $\\sim 1.8 \\times 10^4$ K.  These parameters are of use in attempting to understand the modulation profile in Cet3. For a very small mass-ratio ($q$ = 0.094) the accretion disc nearly fills the Roche lobe of the primary (Warner 1995a) and partial eclipses of the disc are quite probable. The shallow but wide minima in the Cet3 light curve could therefore be eclipse features. The profile could alternatively be a combination of reflection effects from the secondary and aspect effects of the disc and bright spot. However, a small mass ratio $q$ and high $\\dot{M}$ usually result in perturbation of the accretion disc into an elliptical form with consequent precession (Warner 1995a). The AM CVn stars with high $\\dot{M}$ show `superhumps' with periods a few percent longer than $P_{orb}$, which result from the tidal stresses in a precessing disc (Patterson et al.~2001). It is probable, therefore, that the modulation in Cet3 is a permanent superhump, in which case $P_{orb}$ will be slightly shorter than the period we have measured. Alternatively, Cet3 may be a polar like V407 Vul (the absence of two distinct periods disfavours an intermediate polar interpretation). It should be possible to distinguish between these various interpretations when high time resolution spectroscopy is available."
    },
    "0201/astro-ph0201122_arXiv.txt": {
        "abstract": "{ We present highly time resolved circular-polarization and flux spectra of the magnetic white dwarf \\lp\\ taken with the VLT UT1  in order to test the hypothesis that \\lp\\ is a fast rotator with a period of the order of seconds to minutes. Due to low time resolution of former observations this might have been overlooked -- leading to the conclusion that \\lp\\ has a rotational period of over 100 years. The optical spectrum exhibits one   prominent absorption feature with minima at about 4500, 4950, and 5350\\,\\AA, which are most likely C$_2$ Swan-bands shifted by about 180\\,\\AA\\ in a magnetic field between 50\\,MG  and 200\\,MG. At the position of the absorption structures the  degree of  circular polarization varies between -1\\%\\ and +1\\%, whereas it amounts to +8 to +10\\%\\ in the blue and red continuum. With this very high degree of polarization \\lp\\  is very well suited to  a search for short time variations, since a variation of several percent in the polarization can be expected for a magnetic field oblique to the rotational axis. From our analysis we conclude that variations on time scales from 50 to 2500 seconds must have amplitudes $\\lappr 0.7\\%$ in the continuum and $\\lappr 2\\%$ in the strongest absorption feature at 4950\\,\\AA. While no short-term variations could be found a careful comparison of our polarization data of \\lp\\ with those in the  literatures indicates significant variations on time scales of decades with a possible period of about 24-28 years. ",
        "introduction": "\\lp\\ belongs to a  small group of magnetic white dwarfs which were suspected to be polarimetrically constant (West \\cite{west89}; Schmidt \\&\\  Norsworthy \\cite{schmidt91}) implying rotational periods of more than  100 years. Alternatively,  the non-detection of variations in the polarization can also mean a rotationally symmetric magnetic field geometry or a rotational period too short (P$<$10 minutes, Schmidt \\& Norsworthy \\cite{schmidt91}) to be time-resolved as yet. The shortest  period of 12 minutes found in a (magnetic) white dwarf was measured in RE J0317$-$853 (Barstow et al. \\cite{barstow95}; Burleigh et al. \\cite{burleigh99}) meaning either that angular momentum is conserved to a higher degree or that  the star is the product of merging. Segretain et al. (\\cite{segretain97}) predicted rotational velocities of about 1000\\,km/sec ($P\\approx\\,40$ sec) from their models of merging white dwarfs.  If  angular momentum is completely conserved during stellar evolution, rotational periods just above the break-up limit of a few seconds are even possible in the case of single white dwarfs. If any of the strongly magnetic white dwarfs really turns out to be an extremely fast rotator it  may  generate a significant anisotropic moment of inertia and thereby gravitational radiation measurable with space interferometers (Heyl \\cite{heyl00}).  Since such a fast rotation  cannot be excluded from previous observations we obtained highly  time-resolved flux and circular-polarization spectra of \\lp, continuing our search for short rotational periods in white dwarf. Friedrich \\& Jordan (\\cite{friedrich01}) have  started such a search with a broad-band photometric study of the linear polarization in the famous magnetic white dwarf \\grw .  So far, no indications for fast rotation could be found, which  means that the angular momentum has been almost completely lost during stellar evolution. Recently, long-term  variations in the polarization of  \\gd\\ and G$\\,$240-7 have been observed, from which a rotational period of about 100 years can be deduced (Berdyugin \\& Piirola \\cite{berdyugin99}). The puzzling slow rotation of most white dwarf stars is usually explained by magnetic braking of the stellar core during the process of bipolar outflows that may produce the observed bipolar planetary nebulae (Blackman et al. 2001).   Even without a magnetic field  a third dredge-up can very efficiently transport angular momentum from the core of an AGB star to the envelope so that the final white dwarf essentially stops rotating ($v_{\\rm rot}=10^{-3}$ km/sec, Driebe \\&\\ Bl\\\"ocker \\cite{bloecker01}).   With its very strong  wavelength dependence of the degree of circular polarization (-1\\% to +10\\%) \\lp\\ is best suited for  such a study. If  \\lp\\ had a magnetic field not exactly aligned to its rotational axis, a variation of several percent can be expected. In the extreme case of a rotational axis perpendicular to both the observer and the magnetic field axis the polarization in the continuum could vary between 0\\%\\ and 20\\%\\ in order to account for the mean observed value in the continuum of about 10\\%. The only shortcoming of observing this particular object is that no reliable quantum mechanical calculations exist as yet for the C$_2$  molecule features seen in \\lp, but simple estimations explain the shifts of the Swan bands by 500\\AA\\ with magnetic fields strengths between 50\\,MG (Bues \\cite{bues99}) and 200\\,MG (Liebert et al. \\cite{liebert78};  Schmidt et al.\\cite{schmidt95}); therefore  flux and polarization cannot be compared to theoretical models  in order to measure the  Doppler broadening.   \\begin{figure}[htbp] \\includegraphics[width=0.5\\textwidth]{fluxlpall.ps} \\caption[]{Phase resolved and Savitzki-Golay filtered flux spectra of \\lp\\ of the July 3-4, 2000, observing run. The spectra are shifted proportional to the time past since the first observation  (total time: 4800 seconds). The zero point of the ordinate corresponds to the first (upper) exposure. For all observations a line is drawn to indicate the zero level.} \\label{fluxlpall} \\end{figure}  \\begin{figure}[htbp] \\includegraphics[width=0.5\\textwidth]{savpollpall.ps} \\caption[]{Circular polarization of \\lp\\ for 26 phases of the second observation period covering a total of 4600 seconds. The noise was reduced with a Savitzki-Golay filter of  300\\,\\AA\\ width  and a 4th degree polynomial. The zero points of the degrees of polarization for all observations are shown as horizontal lines. Data taken with more than 100 seconds exposure time are marked with thick lines. } \\label{savpollpall} \\end{figure}  \\begin{figure}[htbp] \\includegraphics[width=0.5\\textwidth]{fluxm.ps} \\caption[]{Mean flux spectra  of \\lp\\ and of the unpolarized comparison star.} \\label{fluxm} \\end{figure}   \\begin{figure}[htbp] \\includegraphics[width=0.5\\textwidth]{savpollpallsek.ps} \\caption[]{Comparison of the  mean of our VLT measurements of circular polarization  (curve) to data from the literature: February 22, 1977 (triangles),  and February 23, 1977 (diamonds), taken by  Liebert et al. (\\cite{liebert78}), May 7$-$8, 1994, (plus signs) observed by  Schmidt et al. (\\cite{schmidt95}). For a clearer distinction these measurements were not connected. The  broad-band data by West (\\cite{west89}), marked  by error bars, were taken between January, 1986, and September 1988, the measurement in the red by Robert \\&\\ Moffat (\\cite{robertmoffat89}, February 23, 1987, 'x' symbol)  and  by  Beuermann \\&\\ Reinsch (\\cite{beuermann01}; dashed line) taken on February 4, 2000. } \\label{secular} \\end{figure} ",
        "conclusions": "Up to now, no changes of the degree of circular polarization of \\lp\\ have been reported, resulting in the conclusion that the rotational period exceeds 100 years. As an alternative to such extremely slow rotation we have tested highly time-resolved  circular-polarization  spectra obtained with the VLT for variations on time scales of 50 up to 2500 seconds. In the extreme case of a rotational axis perpendicular to both the observer and the magnetic field axis the polarization in the continuum could vary by about $\\pm 10$\\%. However, we found  no evidence for  rapid variations exceeding 0.7\\% in the continuum or 2\\%\\ at the wavelengths of the absorption features. Smaller variations showing up in our observations are compatible with the assumption that they are produced by statistical noise.  The result is also compatible with fast polarimetric $R$-band photometry (5700-7400\\,\\AA) by Beuermann \\&\\ Reinsch (\\cite{beuermann01}) who found no significant periodicities between 4 seconds and 1.5 hours  in their complementary study of \\lp\\ which extends to longer continuum wavelengths and looks for even shorter rotational periods.  Rapid variations cannot completely be excluded  on a sub-percent level. However, since Friedrich \\&\\ Jordan (\\cite{friedrich01}) also could not measure any fast variations in the case of \\grw\\ fast-rotation as a general scenario for  magnetic white dwarfs with apparently constant flux and polarization  becomes more and more improbable.  The quality of  our mean circular-polarization spectrum considerably exceeds any previous data of \\lp. We compared it to all other measurements of circular polarization published in the literature and found that both the first data  of this object taken in 1977 by  Liebert et al. (\\cite{liebert78}) and data obtained in 1994 by Schmidt et al. (\\cite{schmidt95}) were discrepant from our newer data in the continuum at $\\lambda > 5300$\\AA.  Together with broad band data by West (\\cite{west89}) from 1986-1988, Robert \\&\\ Moffat (\\cite{robertmoffat89}) from 1987, and   Beuermann \\&\\ Reinsch (\\cite{beuermann01}) we conclude that  the degree of circular polarization actually changes systematically and that there is a slight indication for a  period between 24 and 28 years. Taking this period at face value the polarization at $5300-6300$\\AA\\ should increase for the next 9-13 years. Since the question how much of the angular momentum is lost during stellar evolution is of  crucial importance it may be worth monitoring \\lp\\ (and other magnetic white dwarfs) with large telescopes and long-exposure polarization spectra.  For this purpose one high signal-to-noise circular-polarization spectrum (of one or two hours exposure time) with the VLT every year  would be sufficient."
    },
    "0201/astro-ph0201408_arXiv.txt": {
        "abstract": "At this conference many interesting talks were presented on the plausible existence of Quark Stars. Other talks dealt with  the exotic new phases  of quark matter at very high density. Here, I show how combining these two elements might offer a new way of tackling the Gamma Ray Burst puzzle. \\vskip0.5cm ",
        "introduction": "It is widely accepted that the most conventional interpretation of the observed Gamma-ray bursts (GRBs) result from the conversion of the kinetic energy of ultra-relativistic particles to radiation in an optically thin region \\cite{kouveliotou95, kulkarni99,piran99a,piran99b}. The particles being accelerated by a fireball mechanism (or explosion of radiation) taking place near an unknown central engine \\cite{goodman86,shemi90,paczynski90}. The first challenge is to conceive of circumstances that would create a sufficiently energetic fireball. In the model presented in this talk, the approach is to make use of intrinsic properties of quark stars (where exotic phases of quark matter come into play) to account for the fireball.  Quark matter at very high density  is expected to behave as a color superconductor (see Figure 1). A novel feature of such a phase (in the 2-flavor case; hereafter 2SC) is the generation of  glueball like particles (hadrons made of gluons) which as demonstrated in \\cite{ouyed01a} immediately decay into photons. If  2SC  sets in at the surface of a quark star the glueball decay becomes a natural mechanism for a fireball generation; a mechanism which is fundamentally different from models where the fireball is generated via a collapse \\cite{blandford77,ruffert99,janka99} or conversion (of neutron star to quark star \\cite{olinto87,cheng96,bombaci00}) processes. I will then show how and why quark stars might constitute new candidates for GRB inner engines \\cite{ouyed01b}. ",
        "conclusions": "\\label{discuss}   \\subsection{Existence and formation of HQSs}  In the last few years, thanks to the large amount of fresh observational data collected by the new generation of X-ray and $\\gamma$-ray satellites, new observations suggest that the compact objects associated with the X-ray pulsars, the X-ray bursters, particularly the SAX J1808.4-3658, are good quark stars candidates (see Bombaci in this volume and \\cite{li99}). If one assumes that these plausible quark stars form via the ``standard'' supernova mechanism or by conversion of neutron stars then the  two regimes (Heavy and Light stars) discussed in our model are difficult to account for. It has been argued however that quark stars formation mechanisms may be numerous and ``exotic'' (early  discussions can be found in \\cite{alpar87,glendenning97}). In the case of 4U 1728-34 (where a mass of less than $1.0M_{\\odot}$ was derived; \\cite{bombaci00}), it  seems that accretion-induced collapse of white dwarfs is a favored formation mechanism. If the quark star formed via the direct conversion mechanism then it required too much mass (at least $\\sim 0.8M_{\\odot}$ to be ejected during the conversion). Other formation scenarios are discussed in \\cite{hong01,ouyed01c}.   \\subsection{Neutrino cooling and HQSs} \\label{cooling}  If it turns out that neutrino cooling is still very efficient in the 2SC phase, one must consider the scenario where the entire HQS enters the 2SC phase (for comparison of cooling paths between quark stars and neutron stars and the plausible effects of 2SC on cooling we refer the interested reader to \\cite{schaab00,blaschke00,blaschke01}). Here, the 2SC/LGB/photon process (the fireball) occurs only once and inside the entire star. Furthermore, one must involve more complicated physics (such as that of the crust) to account for the episodic emissions so crucial to any model of GRBs. It is not clear at the moment how to achieve this goal and is left as an avenue for future research.  \\subsection{2SC-II stars}  The 2SC/LGB/photon process might proceed until one is left with an object made entirely of 2SC. We name such objects  {\\it 2SC-II} stars\\footnote{The ``II'' in 2SC is a simple reminder of the final state of the star, namely the 2SC with only 5 gluons.} which are still bound by strong interactions (their density is constant $\\sim \\rho_{HQS}$). 2SC-II stars carry an Iron/Nickel crust left over from the GRB phase. The crust mass range is $0 <M_{2SC,crust}<5\\times 10^{-5}M_{\\odot}$ depending on the efficiency of crust extraction/ejection during the GRB phase.   BATSE  observes on average one burst per day. This corresponds, with the simplest model - assuming no cosmic evolution of the rate - to about once per million years in a galaxy \\cite{piran99a}. In the Milky way we thus expect up to $10^{5}$ of 2SC-II stars. Nevertheless, they are tiny enough ($M\\le 10^{-2}M_{\\odot}$, $R\\le 1$ km) to be difficult to detect.    \\subsection{The Mass$-$Radius Plane}  Take observed GRBs with known energies and total duration. From the  burst total energy  $E_{GRB} \\simeq 0.4 Mc^2$ we derive the mass while the burst total duration ($t_{tot}$) gives us the radius (using eq(\\ref{tmassive}) with $T_{c}\\simeq 10$ MeV). In Figure 3 we plot the resulting Mass$-$Radius. The solid curve shows the $M_{HQS}={4\\pi\\over 3}\\rho_{HQS}R_{HQS}^3$ equation which is a reasonable approximation for HQSs. While the GRB data set used is limited nevertheless it seems to support the idea that extremely compact objects ($M\\propto R^3$) are behind GRBs activity within our model.   \\begin{figure}[t!] \\begin{center} \\epsfig{file=ouyedf3.eps,height=3.5in} \\caption{The Mass$-$Radius plane derived in our model using few existing GRBs with known energies and total duration.  The solid curve shows the $M_{HQS}={4\\pi\\over 3}\\rho_{HQS}R_{HQS}^3$ equation for $\\rho_{HQS} \\simeq 9\\rho_{N}$.} \\label{figure3} \\end{center} \\end{figure}"
    },
    "0201/astro-ph0201314_arXiv.txt": {
        "abstract": " ",
        "introduction": "The total number of quark-lepton families is not known for sure. Experimentally, three generations have been found at present. Theoretically, the number of generations could be larger. In superstring-inspired particle models the number of generations is defined by some topological characteristics of the manifold of compactified additional dimensions, and any number of generations is a priori possible \\cite{Green}.  Another important prediction of almost all realistic superstring-inspired models is the existence of at least one additional $U(1)$ gauge group in the low energy limit of the theory. Recently \\cite{Fargion99} it was suggested to ascribe such a new $U(1)$ gauge group to an additional, fourth, fermion generation only. In this case the new gauge group can remain unbroken. Such an unbroken $U(1)$ gauge group implies the existence of a strictly conserved charge, which, in turn, accounts for the stability of the lightest particle of the 4th generation, and forbids any mixing with the other (usual) three generations. It will be assumed that the lightest particle of the 4th generation is its neutrino.  The direct search for new generation fermions on accelerators leads to a lower bound on their masses in the range 50-100 GeV. The mass of a 4th generation neutrino $m_\\nu $ is restricted by the measurement of the width of the Z-boson to $m_\\nu >m_Z/2\\approx 45$ GeV if it is a (quasi)stable Dirac neutrino (for a Majorana one the restriction is slightly lower; for an unstable one the lower limit is about 90 GeV). Another possibility to search for a fourth generation neutrino, with mass $m_\\nu >m_Z/2 \\approx 45$ GeV, at accelerators was suggested in \\cite{Fargion99}, \\cite{Fargion96}. A detailed analysis of the data on the parameters of the Standard Model, accounting for the possible contributions of virtual new generation fermions, allows for the existence of an additional generation if the mass of the new (fourth) neutrino is about 50 GeV \\cite{Maltoni} and if the masses of the other fermions of the new generation exceed 100 GeV.  The existence of new generation fermions in the Universe can lead to many observable astrophysical effects. This makes the appropriate cosmological and astrophysical analysis an important tool for probing the possible existence and properties of such 4th generation fermions. In particular, a new neutrino, being plausibly the lightest of its generation, and thus possibly stable, is of the most interest in such an analysis.  As was found long ago \\cite{Gershtein}, the existence of a heavy Dirac neutrino is compatible with the upper limit on the total density of the Universe if its mass exceeds 2 GeV. Indeed, for a neutrino mass greater than 3 MeV, neutrino annihilation through weak interactions reduces their cosmological concentration at freeze out. A larger neutrino mass corresponds to a larger annihilation cross section, and therefore to a smaller relic heavy neutrino density   $\\Omega _\\nu = \\rho_\\nu / \\rho_{critical}$ (see Fig.2 below). Moreover, if $m_\\nu = m_Z/2$ the Z-boson annihilation resonance leads to a dip in the relic density. For instance, if $m_\\nu =50$ GeV, close to the Z-boson resonance dip, $\\Omega _\\nu \\approx 10^{-4}$. Such a rare population of neutrinos does not play any significant dynamical cosmological role as a Cold Dark Matter (CDM) contributor. Nevertheless a more refined astrophysical analysis \\cite{Fargion99}, \\cite{Golubkov} showed that the effects of such rare, massive neutrinos can be accessible to some experimental searches and/or astrophysical observations. In the present paper, we shall mainly consider the mass range $ 45 < m_\\nu < 90$ GeV, corresponding to such a sparse population of heavy neutrinos.  Heavy neutrinos, as any form of CDM, must be concentrated in galaxies. The ratio of galactic neutrino density to the mean cosmological density is a model-dependent parameter (denoted here as $\\xi $) which is strongly sensitive to the details of galactic halo formation. In this work we shall assume that the concentration factor $\\xi $ is the same for massive neutrinos and for the dominant contributor to CDM. It is usual to estimate $\\xi $ by taking the ratio between a local (i.e, in the vicinity of the Solar system) density of CDM equal to $ 0.3 \\,\\rm{GeV/cm^{-3}} $, and a mean cosmological density corresponding to $\\Omega _{CDM}=0.3$. This leads to the ``standard'' estimate: $\\xi = 2 \\times 10^{5}$. In most of this paper, we shall assume this standard value (except, when explicitly said otherwize).  In previous papers \\cite{Fargion99}, \\cite{Golubkov} we already analyzed several astrophysical consequences of such a sparse galactic population of heavy neutrinos. In particular we studied \\cite{monochromatic} the (ordinary) neutrino fluxes emitted by the annihilation of heavy neutrino-antineutrino pairs accumulated in the core of the Earth. [Such an accumulation takes place for many types of weakly interacting massive particles (WIMP) \\cite{wimp}.]  Recently, it was pointed out \\cite{Bergstrom} that this annihilation flux can be strongly enhanced by the existence of a ``slow'' Solar-system population of WIMP's, trapped in the gravitational field of the Solar system by an initial inelastic interaction with the Sun \\cite{Damour}. The aim of the present paper is to analyze in detail, in the case of a heavy neutrino WIMP, the density of this new ``slow'' Solar-system population, and its enhancement effect on the annihilation flux from the core of the Earth. We shall show that the existence of the slow population qualitatively improves the sensitivity of underground neutrino data to the effects of 4th neutrino annihilation, and makes these data a significant probe of the existence of a 4th neutrino. ",
        "conclusions": "In the present paper we studied the capture by the Earth, and the annihilation in the Earth core, of hypothetical fourth neutrinos (candidate to a sparse sub-dominant component of galactic CDM). We took into account not only the primary ``fast'' population of neutrinos, but also the recently pointed out secondary ``slow'' population \\cite{Damour}.  It was found that the account of the  slow component is crucially important in the considered neutrino mass range, $ 45 < m_\\nu < 90$ GeV. Indeed, the contribution  to capture of the slow component is larger  by up to the two orders of magnitude than the one of the galactic component (Fig.4).  These results suggest the crucial significance of underground experiments (AMANDA, Super-Kamiokande, Baksan and others) for testing the fourth neutrino hypothesis. For example, Ref. \\cite{AMANDA} has derived the constraint on WIMP annihilation in the Earth, from AMANDA data, under some assumptions on the annihilation channels, and for a WIMP mass larger than 100 GeV. Making a rough extrapolation of this constraint down to a mass of about 50 GeV one finds a potential sensitivity of the already existing underground neutrino data to the annihilation of a 4th neutrino in the Earth, for almost all the considered interval of neutrino mass. Note, that the presence of the slow component of a 4th neutrino plays a crucial role in this potential sensitivity. Of course, this example can serve only as an illustration, since special analysis of the data in the framework of the hypothesis of a 4th neutrino is needed.  The possibility to distinguish, in underground neutrino experiments,  the contribution to annihilation effects of a sparse component of 4th neutrino from the contribution of other WIMPs (presumably dominating in the galactic CDM), results from the combination of several factors. The neutrino capture in the Earth is facilitated by the relatively large weak interaction cross section of a massive neutrino and by the kinematic enhancement of neutrino momentum losses in collisions with iron nuclei. The neutrino annihilation effects in the Earth are strengthened by the relatively large neutrino weak annihilation cross section near $Z$-boson resonance (which is further strongly enhanced by the Coulomb like effect of the new long range interaction), and by the existence of a monochromatic neutrino-antineutrino annihilation channel, specific to a 4th neutrino.   The presence of a slow component qualitatively enhances these factors. The slow component increases by up to 50\\% the number density of 4th neutrinos near the Earth. Owing to their order-of-magnitude smaller mean velocity, the slow neutrinos are more effectively captured by the Earth (by up to two orders of magnitude) than the galactic ones. In the slow component capture the kinematic peak of iron nuclei capture is spread over the whole considered neutrino mass interval, making it accessible to experimental test. The establishment of kinetic equilibrium between neutrino capture and annihilation in Earth makes the predicted annihilation fluxes insensitive to the details of captured neutrino distribution. As a result, the account of the slow neutrino component makes the hypothesis of a stable massive 4th neutrino accessible to underground neutrino experimental tests even under the most unfavourable astrophysical conditions.  To conclude, our work shows that it is important to analyze existing (and future) underground neutrino data with the view of probing the hypothetical existence of a stable fourth generation neutrino with a mass about 50 GeV. The analysis of the data of MACRO, AMANDA, Kamiokande, Baksan and/or Super-Kamiokande is expected to provide an important probe of (and probably stringent constraints on) this hypothesis, especially in the case where one wishes to explain the DAMA event rate by assuming heavy neutrinos. The existence of a slow  4th neutrino component is crucial in such an analysis because it strongly enhances the underground neutrino flux expected from 4th neutrino annihilation in Earth. This analysis can be viewed as a modest step towards the study of heterotic string phenomenology which generically leads to the prediction of an additional $U(1)$, which, in turn, provides a motivation for considering a {\\it stable} 4th generation neutrino."
    },
    "0201/astro-ph0201064_arXiv.txt": {
        "abstract": "The recently discovered correlation between black hole mass and stellar velocity dispersion provides a new method to determine the masses of black holes in active galaxies.  We have obtained optical spectra of Markarian 501, a nearby $\\gamma$-ray blazar with emission extending to TeV energies.  The stellar velocity dispersion of the host galaxy, measured from the calcium triplet lines in a $2\\arcsec\\times3\\farcs7$ aperture, is $372 \\pm 18$ \\kms.  If Mrk 501 follows the \\msigma\\ correlation defined for local galaxies, then its central black hole has a mass of $(0.9-3.4) \\times 10^9$ \\msun. This is significantly larger than some previous estimates for the central mass in Mrk 501 that have been based on models for its nonthermal emission.  The host galaxy luminosity implies a black hole of $\\sim6\\times10^8$ \\msun, but this is not in severe conflict with the mass derived from \\sigmastar\\ because the $\\mbh-\\lbul$ correlation has a large intrinsic scatter. Using the emission-line luminosity to estimate the bolometric luminosity of the central engine, we find that Mrk 501 radiates at an extremely sub-Eddington level of $L/\\ledd \\approx 10^{-4}$. Further applications of the \\msigma\\ relation to radio-loud active galactic nuclei may be useful for interpreting unified models and understanding the relationship between radio galaxies and BL Lac objects. ",
        "introduction": "The tight correlation recently discovered between stellar velocity dispersion and black hole mass in nearby galaxy bulges \\citep{fm00, geb00a} has become the key to our understanding of black hole demographics, as well as a new tool for probing the evolution of galaxies and quasars.  An equally important aspect of the \\msigma\\ correlation is its predictive power.  While dynamical mass measurements are observationally difficult and only possible for a limited number of galaxies, the \\msigma\\ relation makes it possible to obtain a black hole mass estimate good to $\\sim40\\%$ accuracy or better from a single measurement of bulge velocity dispersion.  Having such a straightforward method to estimate black hole masses is a tremendous boon to studies of active galactic nuclei (AGNs), because \\mbh\\ is a fundamental parameter affecting the energetics and emission properties of AGNs.  \\citet{geb00b} and \\citet{fer01} have recently shown that black hole masses derived for Seyfert galaxies via the \\msigma\\ relation are consistent with masses determined by reverberation mapping, providing added confidence that AGNs do follow the same correlation as inactive galaxies.  For some classes of AGNs, using the correlations between \\mbh\\ and host galaxy properties may be the \\emph{only} reliable way to estimate the central masses.  BL Lac objects fall in this category, since more direct methods for determining \\mbh\\ (stellar dynamics or reverberation mapping) cannot be applied.  Markarian 501 is one of the nearest known BL Lac objects.  Its redshift is $z = 0.0337$ \\citep{ulr75}, corresponding to a distance of 144 Mpc for $H_0 = 70$ km s\\per\\ Mpc\\per.  The host of Mrk 501, also known as UGC 10599, is a giant elliptical that appears morphologically normal in ground-based images \\citep{hic82, amc91, sfk93, wsy96, nil99}.  Stellar absorption lines are visible in the nuclear spectrum, along with weak emission lines (Ulrich \\etal\\ 1975; Stickel \\etal 1993).  Mrk 501 is also a $\\gamma$-ray source, and is one of the few extragalactic objects from which TeV emission has been detected \\citep{qui96, bra97}.  In this \\emph{Letter}, we present new optical spectra of Mrk 501.  The \\ion{Ca}{2} triplet lines are clearly detected, albeit substantially diluted by the nonthermal continuum.  The stellar velocity dispersion measured from these lines is $372 \\pm 18$ \\kms.  Using this result in conjunction with the \\msigma\\ correlation, we derive an estimate of the mass of the black hole in Mrk 501. ",
        "conclusions": "The correlation between black hole mass and galaxy bulge luminosity can also be used to obtain black hole mass estimates.  The Mrk 501 host galaxy has a total $B$ magnitude of 14.4 (Stickel \\etal\\ 1993). Correcting for Galactic extinction of $A_B=0.084$ mag \\citep{sfd98} and applying a $K$-correction of 0.16 mag \\citep{pen76}, the galaxy has $M_B = -21.6$ mag for $H_0=70$ km s\\per\\ Mpc\\per.  From the fit to the $\\mbh - L_{\\mathrm{bulge}}$ correlation given by \\citet{kg01}, the expected mass is $\\mbh = 6.1\\times10^8$ \\msun.  While this value is outside our $1\\sigma$ uncertainty range on \\mbh\\ from the \\msigma\\ relation, the scatter in the $\\mbh - L_{\\mathrm{bulge}}$ correlation is more than an order of magnitude in \\mbh\\ at fixed $L_{\\mathrm{bulge}}$ \\citep{kg01}.  Also, the host galaxy luminosity is rather uncertain; the values compiled from the literature by \\citet{nil99} show a range of $\\sim0.5$ mag among different authors. (We note that for the adopted luminosity, the galaxy lies very close to the mean Faber-Jackson relation for nearby ellipticals.)  Thus, the level of disagreement between these two \\mbh\\ estimates is not a cause for concern; M87 is an outlier by a similar amount in the $\\mbh-\\lbul$ correlation.  The \\msigma\\ relation is a much more reliable predictor of \\mbh\\ because its intrinsic scatter is much smaller, less than $40\\%$ \\citep{geb00a} and possibly near zero \\citep{fm00}.  Some previous estimates of the black hole mass in Mrk 501 have been obtained by examination of the spectrum and variability of its nonthermal emission.  For example, \\citet{fxb99} derive \\mbh\\ by combining the variability timescale (from which they estimate the physical size of the emitting region) with the assumption that the $\\gamma$-rays in blazars are produced at $\\sim200$ Schwarzschild radii from the black hole.  For Mrk 501 and seven other $\\gamma$-ray blazars, they find central masses in the range $(1-7)\\times10^7$ \\msun.  \\citet{rm00} interpret the periodic behavior observed in X-ray and $\\gamma$-ray light curves as evidence for a binary black hole in Mrk 501, and propose a geometric model for a jet originating from the less massive black hole to estimate $\\mbh \\lesssim 10^8$ \\msun\\ and $(4-42)\\times10^6$ \\msun\\ for the two components of the binary. Recently, \\citet{wxw01} have devised a method to estimate \\mbh\\ in blazars from the peak luminosity and peak frequency.  Using data for a large sample (but not including Mrk 501), they show that their method implies masses of $\\sim10^{9-11}$ \\msun\\ for low-frequency peaked blazars, and $10^{5-8}$ \\msun\\ for high-frequency peaked blazars.  However, these methods have not been calibrated for galaxies whose black hole masses have been measured dynamically, so it is difficult to assess their accuracy.  Black hole masses below $10^8$ \\msun\\ in BL Lac objects would be difficult to reconcile with the \\mbh--\\lbul\\ correlation, since \\hst\\ imaging has conclusively demonstrated that most BL Lac hosts are luminous ellipticals \\citep{urr00}.  In addition, masses of $\\leq10^8$ \\msun\\ would pose a problem for unified models of radio-loud AGNs, in which BL Lacs are interpreted as radio galaxies oriented with the radio jet along our line of sight \\citep[e.g.,][]{up95}.  Nearby FR I radio galaxies, which would presumably appear as BL Lac objects if viewed pole-on, have black holes with masses in the range $\\sim(0.3-3)\\times10^9$ \\msun\\ \\citep[e.g.,][]{har94,ffj96}.  Thus, for Mrk 501, the central mass derived from either $\\sigmastar$ or $L_{\\mathrm{bulge}}$ is consistent with the range of masses expected from unified schemes.  The Eddington ratio of the active nucleus in Mrk 501 can be determined if its bolometric luminosity is known.  Only indirect estimates of \\lbol\\ are possible, since the observed nonthermal emission is highly beamed.  One way to estimate \\lbol\\ is by comparison with FR I radio galaxies, the unbeamed counterparts of BL Lac objects.  Bolometric luminosities for four nearby FR I radio galaxies (M84, M87, NGC 6251, and NGC 4261) are available from \\citet{ho99}, and nuclear emission-line measurements for these objects are listed by \\citet{so81} and \\citet{hfs97}.  For these four galaxies, the ratio $L$(\\hal+[\\ion{N}{2}])/\\lbol\\ has a mean value of $4.4\\times10^{-3}$ with a scatter of only $\\sim25\\%$ over an order of magnitude range in \\lbol.  Mrk 501 has $L$(\\hal+[\\ion{N}{2}]) $=1.1\\times10^{41}$ erg s\\per\\ (Stickel \\etal\\ 1993).  If its $L$(\\hal+[\\ion{N}{2}])/\\lbol\\ ratio is similar to these local FR I galaxies, as might be expected from unified models, then it has $\\lbol\\approx2.4\\times10^{43}$ erg s\\per\\ and $L/\\ledd \\approx 10^{-4}$.  Mrk 501 is evidently an extremely sub-Eddington accretor; this strengthens previous indications that nearby TeV-emitting blazars are intrinsically weak AGNs \\citep[e.g.,][]{cm99}.  The correlations between \\mbh\\ and host galaxy properties make it possible to consider new tests of AGN unification scenarios, by examining the distributions of black hole masses and Eddington ratios between different classes of AGNs.  Such tests have been performed using the host galaxy luminosity to estimate \\mbh\\ \\citep{tre01}, but a comparison based on stellar velocity dispersion would provide improved constraints on the central masses, and new probes of the demographics of AGN classes such as FR I and FR II radio galaxies and blazars.  Also, it would be worthwhile to test whether the black hole masses differ systematically between high- and low-frequency peaked blazars, as advocated by Wang \\etal\\ (2001). It should be possible to measure stellar velocity dispersions in several other low-redshift BL Lac objects.  At recession velocities beyond 12,000 \\kms, the Ca $\\lambda8542$ line will be redshifted into telluric H$_2$O absorption bands and difficult to measure accurately, but the \\mgb\\ spectral region will still be observable.  We are beginning a spectroscopic survey of additional nearby BL Lac objects in order to address these questions."
    },
    "0201/astro-ph0201252_arXiv.txt": {
        "abstract": "We report the detection of an H$\\alpha$ emission-line structure in the upstream side of the planetary nebula NGC 40, which is predicted by numerical simulations, and which is attributed to Rayleigh-Taylor instability. Such a  Rayleigh-Taylor instability is expected to occur at early stages of the interaction process between the interstellar medium (ISM) and a fast moving planetary nebula, as is the case for NGC 40. We resolved the Rayleigh-Taylor instability `tongues', as well as the flatness of the nebula around the  `tongues', which results from the deceleration by the ISM. ",
        "introduction": "Planetary nebulae (PNe) are expanding ionized nebulae expelled from asymptotic giant branch (AGB) stars. The remnant of the AGB star, i.e., its core, ionizes the nebula as it evolves to become a white dwarf. As the gas expands, its density drops and its expansion starts to be influenced by the ISM. If there is a relative motion between the PN and the ISM, the ISM signature on the PN structure will emerge first on the nebular side facing the ISM, i.e., the upstream direction, where the ISM ram pressure on the nebula is the largest. At late times, when the nebula is relatively large and of low density, this type of interaction forms a bow-shaped nebula, e.g., Sh 2-216 (Tweedy, Martos,\\& Noriega-Crespo 1995), KeWe 5 (Kerber et al.\\ 1998), KFR 1 (Rauch et al.\\ 2000), and SB 51 (PN G 357.4-07.2; Beaulieu, Dopita, \\& Freeman 1999), sometimes with a long extension of the bow-shaped region on the down stream direction, e.g. Abell 35 (Jacoby 1981; Hollis et al.\\ 1996). Since the nebula is large and of low surface brightness at this late stage, telescopes with large fields of view should be used to detect these PNe, as was done by Xilouris et al.\\ (1996) and Tweedy \\& Kwitter (1996) in their study of many ancient PNe interacting with the ISM. Due to these and other (e.g., Borkowski, Tsvetanov, \\& Harrington 1993; Tweedy \\& Kwitter 1994;  Muthu, Anandarao, \\& Pottasch 2000; Kerber et al.\\ 2000, 2002) observational studies, and several theoretical works (e.g., Borkowski, Sarazin, \\& Soker 1990; Soker, Borkowski, \\& Sarazin 1991, hereafter SBS; Villaver, Manchado, \\& Garc\\'ia-Segura 2000; Dopita et al.\\ 2000) the basic physics of this stage of PN-ISM interaction is well understood and there are many examples of PNe evolving through this stage.  The PN-ISM interaction process may be influenced by the ISM magnetic field (e.g., Soker \\& Dgani 1997; Soker \\& Zucker 1997) and several types of instabilities (e.g., SBS; Dgani \\& Soker 1994; see review by Dgani 2000). Of particular interest is the Rayleigh-Taylor (RT) instability, which may fragment the outer halo, hence allowing the ISM to stream and interact with nebular material closer to the central star (Dgani \\& Soker 1998). At early stages the RT instability manifests itself as `tongues' located outside, but connected to, the main nebular shell (e.g., SBS; Dopita et al.\\ 2000). This is a general type of structure found in simulations of other dense clouds moving through a low density medium (e.g., Jones, Kang, \\& Tregillis 1994). To our best knowledge, no such RT `tongues' were reported before for PNe interacting with the ISM. The only claim for a RT instability at an early stage of PN-ISM interaction was made by Zucker \\& Soker (1993; hereafter ZS93) for IC 4593. IN IC 4593 the nebula is compressed in the upstream direction, such that its density becomes higher, and there is a structure protruding from the nebula in the upstream direction. ZS93 term the region which protrudes from the nebula a `bump'. However, in IC 4593 the `bump' is blobby, and does not show any large scale internal structure, i.e., no RT tongues are observed.  In the present paper we report the detection of a structure in the upstream side of the interacting PN NGC 40 (PN G120.0+09.8), which closely resembles the structure of RT instability obtained in numerical simulations. We focus on the RT instability feature, which following ZS93 we term a `bump', although its internal structure reveals features not seen in IC 4593, and which are predicted by numerical simulations. It is quite possible that in NGC 40 the interaction process is at an earlier stage than that of IC 4593. Other properties of NGC 40 are thoroughly studied by Meaburn et al.\\ (1996), and hence will not be discussed here. In addition to strengthening the theory of PN-ISM interaction, we hope that this first example of a RT instability at very early stages of PN-ISM interaction will stimulate further observations in searching for such features. ",
        "conclusions": ""
    },
    "0201/astro-ph0201209_arXiv.txt": {
        "abstract": "With the recent advent of large interferometric surveys, we can now probe the physical conditions, dynamics, and  star-forming  properties  of molecular gas in the central kpc of starbursts and active galactic nuclei. We present results from the high-resolution ($\\sim$ 100 pc) interferometric  survey  of molecular gas in the inner kpc of nearby starbursts (Jogee, Kenney, \\& Scoville 2001a) and  the  ongoing  multi-transition survey of cold, warm, and dense molecular gas in a broad range of active and inactive galactic nuclei (Jogee, Baker, Sakamoto, \\& Scoville 2001b). ",
        "introduction": "Using the Owens Valley Radio Observatory (OVRO) in the past five years, a high resolution ($\\sim$ 100 pc) interferometric CO  (J=1-$>$0)  survey of molecular gas  in the inner kpc of eleven circumnuclear starbursts and non-starbursts has been conducted (Jogee 1999; Jogee et al. 2001a) The sample  shows an order of magnitude variation in circumnuclear molecular  gas content and  star formation efficiency (SFE),  defined as the star formation rate per  unit mass of molecular gas. The  circumnuclear starbursts, characterized by a high SFE, include the brightest nearby  starbursts comparable to M82. The survey investigates the physical conditions, dynamics, and  star-forming  properties  of molecular gas in the inner kpc. Selected results are outlined in $\\S$ 1.1-1.2.   \\subsection{The molecular environment}  Molecular gas in the inner  kpc differs markedly from the  the outer disk (see Table 1). In the inner kpc, gravitational instabilities can only be triggered at high gas densities (few 100-1000 M$_{\\tiny \\odot}$ pc$^{-2}$) due to the large Coriolis and pressure forces resulting from the large epicyclic frequency and velocity dispersion. Gravitational instabilities  can promote star formation  by agglomerating  molecular clouds into complexes where they can grow by collision, coalescence, and accretion. In the inner kpc, the growth  timescale  ($t_{\\rm GI}$)  of gravitational instabilities can be  so  short (a few Myrs) that it is comparable to the lifetime  of  OB stars which destroy molecular clouds. One may therefore expect the fraction of gas converted into stars before cloud disruption to be higher in circumnuclear starbursts than in the outer disk. Detailed studies are required to further investigate this possibility. Another special feature of the inner kpc  is the presence of a high pressure, high turbulence molecular ISM. Such an ISM  is believed to favor   the formation of more massive clusters  (e.g., Elmegreen 1993) and may explain why bright super star clusters in non-interacting spirals tend to occur preferentially in the  inner kpc. Table 1 also shows that the prototypical ultra luminous infrared galaxy (ULIRG) Arp 220 looks like a scaled-up  version of the nearby circumnuclear starbursts.  This raises  the possibility that  ULIRGs may be starbursts which have built an extreme  molecular environment (density and linewidths) in the central part of deep  potential well, through major mergers or interactions.   \\begin{table}[t] \\caption{Molecular gas properties in the inner kpc} \\begin{center} \\renewcommand{\\arraystretch}{1.4} \\setlength\\tabcolsep{5pt} \\begin{tabular}{lccc} \\tableline \\vspace{-0.15cm} {Quantities   }   &  {Outer Disk}      & {Inner r=500 pc}      &  {Inner r=500 pc} \\\\ {}   &  { of Sa-Sc}      & {of sample starburst}          &  {of Arp 220}\\\\ \\tableline (1) $M_{\\rm gas-m}$ [$M_{\\tiny \\odot}$] & $\\le$ few $\\times 10^{9}$ & Few $ \\times$ (10$^{8}$-10$^{9}$)   & $ 3 \\times 10^{9}$     \\\\ (2) $M_{\\rm gas}$/$M_{\\rm dyn}$  [\\%] & $<$ 5    & 10 to 30      & 40 to 80      \\\\ (3) SFR [$M_{\\tiny \\odot}$ yr$^{-1}$] & -  & 0.1-11 &  $>$ 100  \\\\ (4) $\\Sigma_{\\rm gas,m}$ [$M_{\\tiny \\odot}$ pc$^{-2}$] & 1-100 & 500-3500    & $ 4 \\times 10^{4}$   \\\\ (5)  $\\sigma$  [km s$^{-1}$] & 6-10 & 10-40 & 90  \\\\ (6) $\\kappa$ [km s$^{-1}$  kpc$^{-1}$] & $<$ 100 & 800-3000   & $>$ 1000   \\\\ (7) $\\Sigma_{\\rm crit}$ [$M_{\\tiny \\odot}$ pc$^{-2}$]& $<$ 10   & 500-1500   & $>$ 2000  \\\\ (8) $t_{\\tiny \\rm GI}$ [Myr]& $>$10    & 0.5-1.5    & $<$ 1   \\\\ (9) $\\lambda_{\\tiny \\rm J}$ [pc& Few $\\times$ (100-1000)    & 100-300   & 90  \\\\ \\tableline \\tableline \\end{tabular} \\end{center} \\label{apptab1b} \\small ~ Rows are : (1) $M_{\\rm gas-m}$,  the molecular gas mass (2) $M_{\\rm gas,m}$/M$_{\\rm dyn}$, the ratio of  molecular gas mass to dynamical mass; (3) SFR, the star formation rate; (4) $\\Sigma_{\\rm gas,m}$ the surface density; (5) $\\sigma$,  the velocity dispersion; (6) $\\Sigma_{\\rm crit}$, the critical density for the onset of gravitational instabilities (7) $\\kappa$,  the epicyclic frequency; (8) $t_{\\tiny \\rm GI}$ = Q/$\\kappa$, the growth timescale of the most unstable wavelength (9) $\\lambda_{\\tiny \\rm J}$, the Jeans length \\normalsize \\end{table}   \\subsection{Molecular gas distribution and triggers of star formation}  The  molecular gas shows a wide variety of  morphologies (Fig.~1) ranging from   relatively axisymmetric annuli or disks (NGC\\,4102, NGC\\,3504, NGC\\,4536,   and  NGC\\,4314), elongated double-peaked  and spiral morphologies (NGC\\,2782, NGC 3351, and NGC 6951)  to extended distributions elongated along the large-scale bar (NGC\\,4569). We find  large gas concentrations inside  the outer inner Lindblad resonance (ILR)  of the bar, consistent with theory (e.g., Combes \\& Gerin 1985). Assuming the  epicycle theory for a weak bar, it is estimated that in the sample galaxies, typically the bar pattern speed is $> 40$--115\\,km~s$^{-1}$~kpc$^{-1}$, the  radius of the outer ILR is  $>$\\,500\\,pc, and the radius of the inner ILR is  $<$\\,300\\,pc. (Jogee et al.  2001a).  \\begin{figure}[h] \\plotfiddle{jogee_fig1.eps}{2.5in}{0}{45}{45}{-150}{-180} \\vspace{2.0in} \\caption[]{ In the SFR/M$_{\\mathrm{H2}}$ vs.~M$_{\\mathrm{H2}}$ plane, the CO  intensity  (contours) is  overlaid on the star formation  (greyscale), as traced by RC  and H$\\alpha$. The dotted line is  the P.A. of the  large-scale stellar bar/oval. The synthesized CO beam is 100--200\\,pc. } \\end{figure}   The starbursts and  non-starbursts have  circumnuclear SFR of  3--11   and  0.1--2\\,M$_\\odot$\\,yr$^{-1}$, respectively. For a given  CO-to-H$_{\\rm 2}$ conversion factor, the starbursts have a larger peak gas surface density $\\Sigma_{\\rm gas-m}$ in the inner 500\\,pc radius than non-starbursts with a similar circumnuclear gas content (Fig.~2a). In the regions of intense star formation, $\\Sigma_{\\rm gas-m}$ remains close to the Toomre (1964) critical density ($\\Sigma_{\\rm crit}$) for the onset of gravitational instabilities, despite an order of magnitude variation in $\\Sigma_{\\rm crit}$  (Fig.~3b and e). In the non-starbursts, there are gas-rich regions with no appreciable  star formation, for instance, at the CO peaks in NGC\\,6951 and inside the ring of HII regions in NGC\\,3351 and NGC\\,4314. The gas surface density, although high,  is still sub-critical in regions of inhibited star formation, as illustrated for NGC\\,4314 in Fig.~3e--f.  \\begin{figure}[t] \\plotfiddle{jogee_fig2.eps}{1.0in}{0}{38}{38}{-150}{-160} \\vspace{0.65in} \\caption {\\bf(a) \\rm The azimuthally-averaged  molecular gas surface density and \\bf (b) \\rm the extinction-corrected SFR per unit area. } \\end{figure}  \\begin{figure}[h] \\plotfiddle{jogee_fig3.eps}{2.0in}{0}{55}{55}{-180}{-130} \\vspace{-0.1in} \\caption[]{ \\rm {\\bf (a, d)} CO   (contours) on H$\\alpha$ (greyscale) distributions.  {\\bf (b, e)} $\\Sigma_{\\rm gas-m}$, $\\Sigma_{\\mathrm{crit}}$, and  $\\Sigma_{\\mathrm{shear}}$. {\\bf (c, f)} Toomre Q and shear S parameters (Jogee et al 2001a). Quantities are plotted starting at a radius equal to the CO beam  size ($\\sim$2$''$). See text for details. } \\end{figure} ",
        "conclusions": ""
    },
    "0201/astro-ph0201386_arXiv.txt": {
        "abstract": "The gas temperature in the cores of many clusters of galaxies drops inward by about a factor of three or more within the central 100~kpc radius. The radiative cooling time drops over the same region from 5 or more Gyr down to about $10^8\\yracf$. Although it would seem that cooling has taken place, XMM and Chandra spectra show no evidence for strong mass cooling rates of gas below 1--2 keV. Chandra images show holes coincident with radio lobes and cold fronts indicating that the core regions are complex. The observational situation is reviewed here and ways in which continued cooling may be hidden are discussed, together with the implications for any heat source which balances radiative cooling. ",
        "introduction": "The gas density within the central 100 kpc or so of the centre of most clusters of galaxies is high enough that the radiative cooling time of the gas is less than $10^{10}$~yr. The cooling time drops further at smaller radii, suggesting that in the absence of any balancing heat source much of the gas in the central regions is cooling out of the hot intracluster medium. In order to maintain the pressure required to support the weight of the overlying gas, a slow, subsonic inflow known as a cooling flow develops.  X-ray observations made before Chandra and XMM-Newton were broadly consistent with the cooling flow picture (see 23 for a review and 42 for an opposing view), although several issues remained unresolved. The first issue was the observed X-ray surface brightness profile, which was not as peaked as expected from a homogeneous flow. Instead a multiphase gas was assumed, dropping cold gas over a range of radii. The second was the fate of the cooled gas. At the rates of 100s to more than $1000\\Msunpyracf$ found in some clusters, the central galaxies should be very bright and blue if the cooled gas forms stars with a normal intial-mass-function. In many cases they do have excess blue light indicative of massive star formation [36, 1, 13, 15], but at rates which are a factor of 10 to 100 times lower than the X-ray deduced mass cooling rate. It has been argued [46] that there is no significant sink in terms of cold gas clouds. A third issue involved the shape of the soft X-ray spectrum, which was inconsistent with a simple cooling flow. Absorption intrinsic to the flow was found to be a possible explanation [2, 3]. A final, major, issue was whether the neglect of heating is justified. The effect of gravitational heating as the gas flows was taken into account, but the effects pf any central radio source, which pumps energy into the surrounding gas via jets, together with disturbances due to subclusters plunging into the core every few Gyr were not included due to a lack of quantitative information. Heat flow due to thermal conduction was also generally assumed negligible.  The situation with cluster cooling flows has been clarified over the past year, particularly by the high spatial resolution imaging of Chandra and the high spectral resolution of the XMM-Newton Reflection Grating Spectrometer (RGS). Chandra images show much detail in the cores of clusters, with bubbles from radio sources [41, 24] and cold fronts [40, 57] seen. RGS spectra [50, 54, 38] confirm the presence of a range of temperatures in cooling flow clusters but fail to show evidence for gas cooling below 1--2~keV. Simply put, the data are consistent with gas cooling at a high rate to about one third of the mean temperature beyond 100~kpc but then vanishing.  At about the same time, the evidence for both warm [35,17, 20] and cold [19] molecular gas at the centres of cooling flows clusters has become widespread. In some extreme cases there may be over $10^{11}\\Msunacf$ of cold gas [19]. The presence of dust in these regions is also widespread, as demonstrated by the Balmer decrement in the optical/UV nebulosities commonly seen (e.g. 32, 15), dustlanes, and submm and IR detections [18, 3, 34].  It is therefore possible that more star formation, and in particular cold gas clouds, may be found in and around central cluster galaxies (see also 29 for a discussion of the properties of very cold gas clouds). There has also been the intriguing detection of OVI emission from A2597 with FUSE [47]. Lastly, recent numerical simulations of evolving cluster which include radiative cooling of the gas predict cooling flows (e.g. 48).  At face value the X-ray data tempt many to assume that some form of heating balances cooling and so dismiss cooling flows altogther. That ignores the how, why and what of the heating, which remains unsolved, although several candidates have been identified [56, 6, 16, 10, 11]. Some form of feedback is probably required to prevent all of the gas from being heated up. If feedback does occur we have a good chance to observe how it works, since the region is spatially resolved and optically thin. The process is of wide importance, since it provides the upper mass limit for galaxies (in simulations of the galaxy luminosity function, [39] switch off cooling in massive galaxies).  My own view is to treat it as an intriguing astrophysical puzzle whch can be tackled observationally.  Heating from radio sources and infalling subclusters must occur at some level, but whether it can balance radiative cooling over the required spatial scales to better than a factor of a few is not yet clear. Cooling probably does account for the observed star formation and cold gas clouds. A major remaining issue is whether the mass cooling rates are reduced from the earlier X-ray deduced rates by a factor of a few, ten or a hundred. We may be witnessing a nearby example of the kind of feedback processes common in galaxy formation; in particular, one in which accretion onto the central black hole and the resultant kinetic energy release play a major role. ",
        "conclusions": "The central 100 kpc radius region in most clusters has a radiative cooling time shorter than 5 Gyr (Peres et al 1997) and many have central cooling times of only $10^8\\yracf$. The gas temperature drops by a factor of 3 or more over this radius range. It seems plausible that the temperature drop and short radiative cooling times are related and that the low temperatures are caused by radiative cooling.  It is then a puzzle as to why gas which has cooled by a factor of three, and for which the radiative cooling time has reduced by a factor of ten or more, is not seen to cool further.  There are two obvious solutions, both of which have some difficulties. The first solution is that the gas does cool but either the soft X-ray emission is absorbed or the cooling is non-radiative and due, say, to mixing. The problem of the fate of cooled gas then remains. The second solution is that some heating balances cooling. The problem here is that the heat has to balance cooling over a wide range of radii and a wide range of timescales. Also observations of radio lobes which are a likely source of heat indicate that they coincide with the coolest gas in cluster cores.  The answer may be more complex, with the major temperature drop being due to a combination of in situ radiative cooling and gas introduced from dense cooling subclusters. Heat from the kinetic energy of infalling subclusters, and turbulence, continues to be dissipated throughout the core, reducing the age of any steady central cooling region to only a few Gyr. An intermittent central radio source powered by accretion from the intracluster medium heats and churns up some of the coolest gas at the centre. Radiatively-cooling clumps (possibly metal rich) fall out of the mean slow inflow once their temperature drops to below one third of the outer temperature and rapidly mix with cooler gas clouds closer to the centre. The mixture would have a temperature of about $10^5\\Kacf$ (as in mixing layers; Begelman \\& Fabian 1990), and rapidly lose its thermal energy by UV emission. Most of this would be absorbed by neighbouring cold gas and dust, to be reradiated as optical/UV line emission and infrared dust emission [28]. Massive star formation take place in cooled clumps and spreads dust into the surrounding gas, further enhancing cooling via infrared emission."
    },
    "0201/astro-ph0201453_arXiv.txt": {
        "abstract": "{ We present chemical abundances for four main sequence B stars in the young cluster NGC\\,2004 in the Large Magellanic Cloud (LMC). Apart from H\\,{\\sc ii} regions, unevolved OB-type stars are currently the only accessible source of present-day CNO abundances for the MCs not altered by stellar evolution. Using UVES on the VLT, we obtained spectra of sufficient resolution ($R$\\,=\\,20\\,000) and signal-to-noise (S/N\\,$\\geq$\\,100) to derive abundances for a variety of elements (He, C, N, O, Mg and Si) with NLTE line formation.\\\\ This study doubles the number of main sequence B stars in the LMC with detailed chemical abundances. More importantly and in contrast to previous studies, we find no CNO abundance anomalies brought on by e.g. binary interaction or rotational mixing. Thus, this is the first time that abundances from H\\,{\\sc ii} regions in the LMC can sensibly be cross-checked against those from B stars by excluding evolutionary effects. We confirm the H\\,{\\sc ii}-region CNO abundances to within the errors, in particular the extraordinarily low nitrogen abundance of $\\varepsilon$(N)\\,$\\simeq$\\,7.0. Taken at face value, the nebular carbon abundance is 0.16\\,dex below the B-star value which could be interpreted in terms of interstellar dust depletion. Oxygen abundances from the two sources agree to within 0.03\\,dex.\\\\ In comparison with the Galactic thin disk at MC metallicities, the Magellanic Clouds are clearly nitrogen-poor environments. ",
        "introduction": "Young populous clusters like NGC\\,2004 in the Large Magellanic Cloud (LMC) are ideal objects to study both the current state of chemical evolution of the LMC and the physics of hot stars at a metallicity a factor of two below that of the Sun.\\\\ The study of stellar absorption spectra yields chemical abundances with an accuracy comparable to that achieved from the analysis of nebulae such as H\\,{\\sc ii} regions. However, both sources are subject to physical processes which alter the abundances of certain chemical species in a systematic fashion. In most cases, the {\\em direction} of these systematic effects is known, yet their extent is often uncertain. In particular, the scaling of these effects with metallicity is poorly known. In H\\,{\\sc ii} regions, it is mostly dust formation and depletion onto grains which will result in lower limits on the present-day ({\\em pristine}) abundances. As for B stars, rapid rotation can lead to contamination of the atmospheres with CN-cycled material from the core resulting in He and N enrichment and a corresponding C (and -- to a lesser extent -- O) depletion (Fliegner et al. \\cite{fliegner}, Meynet \\& Maeder \\cite{maeder}, Heger \\& Langer \\cite{heger}). The onset of ``rotational mixing'' is a function of the rotation rate, stellar mass and metallicity. Rotating models of hot stars indicate that it can take place very early on, i.e. whilst on the main sequence (MS). Furthermore, one can expect rotational mixing to be more efficient in metal-poor environments (cf. Maeder \\& Meynet \\cite{mandm}), as metal-deficient stars have smaller radii for a given mass and thus rotate faster for a given angular momentum. \\begin{table} \\caption[]{Chemical abundances of bona-fide MS B stars in the LMC in comparison with H\\,{\\sc ii}-region data. $\\varepsilon$(X)\\,:=\\,$\\log$($n_{\\rm X}/n_{\\rm H}$)+12. Typical errors are 0.2\\,--\\,0.3\\,dex.} \\label{prevB} \\begin{tabular}{lllll} \\hline star & PS 34-16$^{\\mathrm{a}}$ & LH 104-24$^{\\mathrm{a}}$ & NGC\\,1818/D1$^{\\mathrm{b}}$ & H\\,{\\sc ii}$^{\\mathrm{c}}$ \\\\ \\hline $\\varepsilon$(C) & \\ \\ \\ 7.10 & \\ \\ \\ 7.50 & \\ \\ \\ \\ 7.83 & 7.90 \\\\ $\\varepsilon$(N) & \\ \\ \\ 7.50 & \\ \\ \\ 7.70 & \\ \\ \\ \\ 7.59$^{\\mathrm{d}}$ & 6.90\\\\ $\\varepsilon$(O) & \\ \\ \\ 8.40 & \\ \\ \\ 8.50 & \\ \\ \\ \\ 8.46 & 8.40 \\\\ $\\varepsilon$(Mg) & \\ \\ \\ 7.00 & \\ \\ \\ 7.40 & \\ \\ \\ \\ 7.35 & \\ \\ -- \\\\ $\\varepsilon$(Si) & \\ \\ \\ 7.00 & \\ \\ \\ 7.40 & \\ \\ \\ \\ 7.10 & 6.70 \\\\ $\\varepsilon$(Fe) & \\ \\ \\ 7.20 & \\ \\ \\ \\ \\ -- & \\ \\ \\ \\ 7.34 & \\ \\ -- \\\\ \\hline \\end{tabular} \\begin{list}{}{} \\item[$^{\\mathrm{a}}$] from Rolleston et al. (\\cite{rolleston}), $^{\\mathrm{b}}$ from Korn et al. (\\cite{korn}), \\item[$^{\\mathrm{c}}$] from Garnett (\\cite{garnett}), $^{\\mathrm{d}}$ revised in this publication\\\\[-6ex] \\end{list} \\end{table}  \\noindent Among a sample of B stars only the fastest rotators ($v_{\\rm rot}>$\\,200 km/s) are expected to display contaminated CNO abundances (Heger \\& Langer \\cite{heger}, Meynet \\& Maeder \\cite{maeder}). If one is interested in the pristine present-day abundance pattern of these elements, it is therefore sensible to look for and analyse the least evolved objects accessible.\\\\ So far, the faintness of unevolved MC B stars ($m_{\\rm V}>15^{\\rm m}$) has prevented an extensive confrontation of the H\\,{\\sc ii}-region CNO data with that from B dwarfs. Four stars have been studied in detail, three by Rolleston et al. (\\cite{rolleston}) and one by Korn et al. (\\cite{korn}, hereafter Paper I). Of the three stars of Rolleston et al. (\\cite{rolleston}) one (PS 34--144) is classified as a He-weak star which we will therefore disregard in what follows. The CNO abundances of the remaining three MS stars are confronted with the most recent measurements from H\\,{\\sc ii} regions in Table \\ref{prevB}: While the consistency among determinations of O abundances is very high, there is a significant positive offset ($\\simeq$ 0.7 dex) of B-star N abundances with respect to the H\\,{\\sc ii}-region value.\\\\ It is a priori unclear which of the two data sets reflects the pristine LMC N abundance. There are many possible explanations for the peculiar behaviour of N: either the H\\,{\\sc ii}-region value is systematically too low (note that it is --~like all the other H\\,{\\sc ii}-region data~-- exclusively based on 30 Dor and N44C, cf. Garnett \\cite{garnett}), the B-star value systematically too high or we see the imprint of a physical process which alters one of the two data sets. Since dust depletion of nitrogen seems hard to envision (Mathis \\cite{mathis}), it looks as if the B stars are to blame. And indeed, it might be that -- due to the limitations imposed on us by the use of 4m-class telescopes up to 1999 -- choosing the brightest targets compatible with MS colours resulted in preselecting stars that are preferentially rapid rotators as rapidly rotating stars will evolve to higher luminosities (Fliegner et al. \\cite{fliegner}).\\\\ Whatever the scenario, we regard it as worthwhile to push the limit further towards the zero-age MS to see whether the N abundances already found are representative of LMC B stars in general.\\\\ In contrast to the paucity of data on dwarfs, many MC supergiants -- from spectral type O to K -- have been studied in detail (e.g. Haser et al \\cite{haser}: O-type, Paper I: B-type, Venn \\cite{venn2}: A-type, Andrievsky et al. \\cite{andrievsky}: F-type, Hill \\& Spite \\cite{hill}: K-type). Obviously, in trying to understand the CNO pattern of these evolved objects the problem lies in separating all contributing signatures: abundance changes due to dredge-up episodes in the red-giant phase and potentially rotational mixing in earlier phases when the rotation rate was still high. This is particularly true for stars located in the so-called Blue Hertzsprung Gap (between the MS and the region of blue-loop excursions, cf.\\ Fitzpatrick \\& Garmany \\cite{fitzpatrick}), whose evolutionary status is still highly uncertain.\\\\ Like in the MS B stars already discussed, the most significant signature is that of nitrogen: Among SMC A-type supergiants Venn (\\cite{venn2}) found a large scatter in NLTE (non local thermodynamic equilibrium) N abundances (6.9\\,$\\leq$\\,$\\varepsilon$(N)\\,$\\leq$\\,7.7) which is hard to explain in the framework of the first dredge-up alone and argues in favour of rotational mixing in earlier evolutionary phases as the most likely cause. As far as abundances from other elements with less pronounced signatures are concerned, one has to worry about how severely they might be affected by systematic effects when comparing different spectral types: not in all cases the assumptions made (plane-parallel geo\\-metry, stationarity, LTE) are as justified as in the case of dwarfs. For example, the carbon abundance from K supergiants seems to be systematically higher than in the hotter stars. On the whole, the general picture drawn above is, however, supported by all spectral types.\\\\ All in all, rotational mixing could help to resolve the discrepancies outlined above: adding rapid rotation, a fraction of any young stellar population will display modified CNO abundances, widen the MS, populate regions of the HRD which are traversed quickly by non-rotating stars and lead to enhanced scatter in N after the first dredge-up by mimicking a distribution of initial abundances.\\\\ \\begin{table} \\label{log} \\caption{Observation log for the programme stars. Nomenclature and visual magnitudes from Robertson (\\protect\\cite{robertson}). $t_{\\rm exp}$ denotes the expsoure time, m/y the month and year of the observation. The line-free wavelength ranges 4202\\,--\\,4210, 4402\\,--\\,4409, 4493\\,--\\,4500 and 4680\\,--\\,4695\\,\\AA\\ were used to estimate the S/N ratios (1$\\sigma$) of the rebinned spectra. Appended to this table are the two luminosity class {\\sc iii} stars in NGC\\,2004 from Paper I observed with CASPEC.} \\begin{tabular}{lcrrrr} \\hline object in & m$_V$ & $t_{\\rm \\,exp}\\ $ & S/N & m/y & $v \\sin i$ \\\\ NGC\\,2004 & [mag] & [h]               &     &     & [km/s] \\\\ \\hline B18 & 14.8 & 1.0 + 1.0 & 140 & 12/00 & 130 \\\\ C8 & 14.9 & 1.0 & 110 & 11/00 & 100 \\\\ C9 & 15.8 & 1.4 + 1.4 & 100 & 11/00 & 60 \\\\ C16 & 14.7 & 1.0 & 130 & 11/00 & 60 \\\\ D3 & 15.8 & 1.4 + 1.4 & 120 & 12/00 & 70 \\\\ D15 & 15.1 & 1.5 & 120 & 12/00 & 45 \\\\ \\hline B15 & 14.2 & 3.0 + 3.0 & 70 & 12/89 & 25 \\\\ B30 & 13.8 & 3 x 2.0 & 150 & 11/87 & 30 \\\\ \\hline \\end{tabular} \\end{table} ",
        "conclusions": "We have presented detailed chemical abundances for four MS B stars in NGC\\,2004. This study doubles the number of MS B stars in the LMC for which accurate abundances have been derived by means of high-resolution spectroscopy. None of the stars shows any indication of effects due to binary interaction or rotational mixing allowing us to derive what we consider to be the first data set of {\\em unaltered present-day CNO abundances} from LMC stars. \\begin{itemize} \\item We confirm the extraordinarily low LMC nitrogen abundance previously found from H\\,{\\sc ii}-region studies to within 0.1 dex. This value is more than 0.5\\,dex below average values found in the Galactic thin disk at LMC metallicities (Liang et al. \\cite{liang}, cf. their Fig.~10). It implies an enrichment history for the two environments different from one another and -- within the chemical evolution model of Henry et al. (\\cite{henry})~-- a dominance\\linebreak of primary nitrogen production in the LMC {\\em until\\linebreak today}. \\item In our programme stars, carbon is 0.16\\,dex above the nebular value. While this is the direction in which dust depletion would act, the offset found is hardly significant considering the error limits on both sides. More work is needed to reduce both random and systematic errors to draw definitive conclusions on dust depletion fractions. \\item The B-star oxygen abundance (also from B giants, cf. Paper I) is in excellent agreement with the nebular value. In the absence of systematic offsets between the two data sets, little nebular oxygen seems to be tied up in grains at LMC metallicities. \\item The stellar silicon abundance is offset from the nebular one by 0.4\\,dex. This might imply that a major fraction of the interstellar silicon in the LMC is bound in grains. \\end{itemize} This work illustrates the potential of abundance analyses of hot stars based on high-quality spectra and sophisticated input physics. With modern multi-object spectrographs (e.g. GIRAFFE on VLT UT2 going on-line in 2002) we will be able to analyse some 100 slowly rotating B stars in each Galactic and MC cluster like NGC\\,2004. This will not only result in reduced {\\em random errors} (which we will be able to address truly statistically for the first time), but also prove or refute the importance of rotational mixing for this class of stellar objects. In particular, a sequence consisting of Galactic, LMC and SMC B stars could  settle the r\\^{o}le metallicity, rotation and mass loss play in hot-star evolution. As for the H\\,{\\sc ii}-region data, more objects will have to be observed to beat down the errors on that side as well.\\\\ {\\em Systematic errors} --~arguably more important than random ones at the current level of accuracy~-- will have to be addressed by an integral approach: by studying the stellar and gas component of H\\,{\\sc ii} regions in tandem. Both fields of research would undoubtedly profit from such an endeavour."
    },
    "0201/astro-ph0201179_arXiv.txt": {
        "abstract": "The effect of the Hall term on the evolution of the magnetorotational instability (MRI) in weakly ionized accretion disks is investigated using local axisymmetric simulations.  First, we show that the Hall term has important effects on the MRI when the temperature and density in the disk is below a few thousand K and between $10^{13}$ and $10^{18}$~cm$^{-3}$ respectively.  Such conditions can occur in the quiescent phase of dwarf nova disks, or in the inner part (inside 10 -- 100 AU) of protoplanetary disks.  When the Hall term is important, the properties of the MRI are dependent on the direction of the magnetic field with respect to the angular velocity vector {\\boldmath $\\Omega$}. If the disk is threaded by a uniform vertical field oriented in the same sense as {\\boldmath $\\Omega$}, the axisymmetric evolution of the MRI is an exponentially growing two-channel flow without saturation.  When the field is oppositely directed to {\\boldmath $\\Omega$}, however, small scale fluctuations prevent the nonlinear growth of the channel flow and the MRI evolves into MHD turbulence. These results are anticipated from the characteristics of the linear dispersion relation.  In axisymmetry on a field with zero-net flux, the evolution of the MRI is independent of the size of the Hall term relative to the inductive term. The evolution in this case is determined mostly by the effect of ohmic dissipation. ",
        "introduction": "The structure and evolution of accretion disks are largely determined by angular momentum transport processes.  One of the most promising processes is MHD turbulence driven by the magnetorotational instability (MRI) (Balbus \\& Hawley 1991).  In ideal MHD, the growth rate of the MRI is of the order of the orbital frequency $\\Omega$, and the characteristic wavelength of the most unstable mode is $2 \\pi v_{\\rm A}/\\Omega$, where $v_{\\rm A}$ is the Alfv{\\'e}n speed.  The nonlinear regime of the MRI has been well studied in ideal MHD using numerical simulations.  However, in some systems, accretion disks are expected to be only partially ionized, and non-ideal MHD effects, which generally suppress the growth of the MRI, must be considered.  For example, the low temperatures of accretion (protoplanetary) disks around young stellar objects make thermal ionization processes ineffective, so that the abundance of charged particles is very small, and non-ideal MHD effects are important (Gammie 1996; Stone et al. 2000).  Another example is provided by dwarf nova systems in quiescence.  The temperature of the disk in this case can be well below $10^4$ K and the disk is only weakly ionized, so that non-ideal MHD effects are again important (Gammie \\& Menou 1998).  There are three regimes in non-ideal MHD associated with the relative importance of different terms in the generalized Ohm's law (see \\S 2); they are the ambipolar diffusion, ohmic dissipation, and Hall regimes.  Which term dominates depends on the ionization fraction and density of the gas.  Ambipolar diffusion is most important in regions of relatively low density and high ionization (e.g., Reg{\\H o}s 1997).  The linear properties of the MRI in the ambipolar regime have been explored by Blaes \\& Balbus (1994); they find unstable modes exist when the collision frequency of an ion with neutrals is higher than the orbital frequency.  The nonlinear evolution of the MRI in this regime was examined by Hawley \\& Stone (1998) using two fluid simulations.  They found that when the coupling between ions and neutrals is weak, the turbulence in the ionized component of the plasma excited by the MRI does not affect the motion of neutrals very much, thus significant angular momentum transport requires a greater coupling between the ions and neutrals than that required for linear instability. Brandenburg et al. (1995) and Mac Low et al. (1995) also studied the effect of the ambipolar diffusion in the strong coupling limit in a few models.  Ohmic dissipation becomes important when the ionization fraction of the gas is very low.  In this case, a linear analysis (Jin 1996; Sano \\& Miyama 1999) shows that small wavelength perturbations are damped, and the characteristic wavelength of the MRI increases in comparison to the ideal MHD case.  The axisymmetric 2D evolution of the MRI demonstrates that nonlinear saturation can occur due to ohmic dissipation (Sano, Inutsuka, \\& Miyama 1998), even though the corresponding ideal MHD cases show an ever-growing channel flow without saturation (Hawley \\& Balbus 1992).  Fleming, Stone, \\& Hawley (2000) examined the nonlinear evolution using local 3D simulations; they found that dissipation weakens the MHD turbulence.  For significant turbulence and angular momentum transport to occur, a critical value for the magnetic Reynolds number must be exceeded, and this value depends on the field geometry in the disk.  Recently, linear analyses of the MRI in the Hall regime have been presented by Wardle (1999) and Balbus \\& Terquem (2001).  The maximum growth rate and characteristic wavelength of the MRI are strongly modified by the Hall effect.  Most interesting is that the linear properties of the instability depend on the direction of the magnetic field.  This is because the dispersion relation for incompressible Alfv{\\'e}n waves traveling along field lines is quite different in Hall MHD.  In particular, the left- and right-circularly polarized Alfv{\\'e}n waves have different phase velocities, and these two waves interact with the Coriolis force in the disk in different ways.  One of these two waves is commonly referred to as the whistler wave.  The purpose of this paper is to investigate the effect of the Hall term on the nonlinear evolution of the MRI in axisymmetry. When the Hall term is important, ohmic dissipation often cannot be neglected (Balbus \\& Terquem 2001).  Thus, we solve an induction equation that includes both ohmic dissipation and the Hall effect. The plan of this paper is as follows.  We examine when the Hall term becomes important in dwarf nova and protoplanetary disks in \\S 2.  Our numerical method and the initial conditions are described in \\S 3. The results of 2D MHD simulations are presented in \\S 4 for both a uniform and a zero-net flux vertical field.  The application of the results to actual accretion disks is discussed in \\S 5, and our results are summarized in \\S 6. ",
        "conclusions": "\\subsection{Interpretation of the 2D Evolution}  The axisymmetric MRI evolves into a channel flow or MHD turbulence, depending on the parameters $Re_{M0}$ and $X_0$ (provided linearly unstable modes exist at all).  The results summarized in Figure \\ref{fig:fate} reflect the linear properties of the MRI.  From the dispersion relation, the parameter space $X_0 > -4$ can be divided into three regions, as shown by dotted lines in Figure \\ref{fig:fate}.  When $Re_{M0} \\gtrsim 1$ and $X_0 \\ge 0$, the characteristic wavelength of the MRI is proportional to $v_{A} / \\Omega$, and a critical wavelength exists.  For this case, the evolution of the MRI shows an inverse cascade of the magnetic energy, and a two-channel flow emerges without saturation.  This is because the characteristic scale increases as the field strength is amplified by the growth of the MRI.  In the second region ($Re_{M0} \\lesssim 1$), the dispersion relation has a critical wavelength, but the most unstable wavelength is proportional to $\\eta / v_{\\rm A}$, that is inversely proportional to the field strength.  In the third region ($X_0 < 0$), there is no characteristic scale of the MRI.  Then smaller scale fluctuations are always unstable.  Therefore, models in these last two regions show no evidence for an inverse cascade; instead they evolve into MHD turbulence.  Models with $\\beta = 800$ also show the same characteristics as the models plotted in Figure \\ref{fig:fate}, implying these results are independent of the initial field strength.  Estimation of the saturated amplitude of the Maxwell stress, which determines the efficiency of angular momentum transport in actual accretion disks, requires 3D simulations of the MRI.  In 3D simulations that include only the ohmic dissipation, the saturation amplitude of the Maxwell stress is larger if the MRI has an inverse cascade (Sano \\& Inutsuka 2001).  Local 3D simulations of the MRI in Hall MHD are presented in Sano \\& Stone (2002).  \\subsection{Definition of Magnetic Reynolds Number}  In this paper we have defined the magnetic Reynolds number as $Re_{M} = v_{\\rm A}^2 / \\eta \\Omega$, and used $Re_{M0}$ to denote the value of this number in the initial state.  Both a linear analysis of the MRI including ohmic dissipation, as well as local 3D simulations with an initial uniform vertical field reveal that the critical number required to generate significant MHD turbulence and angular momentum transport is $Re_{M0} \\sim 1$ (Sano \\& Inutsuka 2001).  If $Re_{M0} > 1$ the evolution and saturated state of the MRI is little changed from the ideal MHD case, whereas if $Re_{M0} < 1$ the growth rates and amplitude of the saturated state are both significantly reduced compared to the ideal MHD case.  In actual accretion disks, there is a minimum value for $Re_{M}$ that results from the requirement that the critical wavelength of unstable modes in the very resistive regime $\\lambda_{\\rm MRI} \\sim \\eta / v_{\\rm A}$ be less than the disk thickness $H$ (Sano \\& Miyama 1999). This requires $Re_{M0} > v_{\\rm A}/c_s$, which can also be used as a stability criterion (Gammie 1996; Igea \\& Glassgold 1999; Sano et al. 2000).  If the field strength of accretion disks is subthermal and $v_{\\rm A}/c_{s} < Re_{M0} < 1$, the MRI will be present, but the growth rate, saturated amplitude, and angular momentum transport rate will all be reduced compared to the ideal MHD case.  In practice, the largest value for the ratio $v_{\\rm A}/c_s$ that can be reached in an accretion disk is probably unity (otherwise the field will escape the disk via buoyancy).  In this case, the magnetic Reynolds number becomes $Re_{M}' = c_s^2 / \\eta \\Omega$ (Gammie \\& Menou 1998; Fleming et al. 2000), with the critical value measured from 3D simulations about $10^{4}$.  This last form is independent of the magnetic field strength in the disk, and therefore can be measured more easily with observations.  However, even if the observed $Re_{M}'$ is below the critical value $10^4$, the MRI may still operate if $0.01 \\lesssim v_{\\rm A} / c_{s} < 1$."
    },
    "0201/astro-ph0201429_arXiv.txt": {
        "abstract": " ",
        "introduction": "The measurement of CO line emission in high redshift objects has proven to be a fruitful avenue for investigating the properties of distant quasars and galaxies. Molecular observations address interesting issues such as that of star formation in the early universe and how the presence of a central massive engine can affect the interstellar material in its host galaxy. Detections of CO have shown that quasars share some properties with luminous infrared galaxies, both locally and at high redshift. For example, comparison of the molecular, infrared and optical properties of the Cloverleaf quasar and the infrared galaxy IRAS F10214+4724 (known to harbor a buried quasar seen in polarized light) has demonstrated that these two objects are nearly identical, except in the optical range where the differences can probably be attributed to obscuration/orientation effects (Barvainis et al.\\ 1995). Such findings lend support to theories unifying luminous infrared galaxies and quasars via orientation effects, with high redshift infrared galaxies being the luminous counterparts of local Seyfert 2's. However, on physical grounds, it also seems likely that IR-selected galaxies and UV-selected quasars may differ in their stage of evolution. There are currently about 15 well-documented detections of molecular gas at high redshift (e.g., Combes 2001), of which at least 9 are gravitationally lensed systems.  The advantages of using an intervening ``gravitational telescope'' to boost the fluxes are obvious, with estimated magnification factors of up to 100 in the optical. Moreover, differential gravitational effects provide an elegant tool to probe the size and structure of the molecular material within the quasar.  For example, a point-like emitting region (rest-frame UV and optical continua from the inner accretion zone), and an extended dusty molecular region (the ``torus'') in the quasar will produce, after gravitational effects from the intervening lens, images with different morphologies.  Molecular line profiles, reflecting intrinsic geometrical and kinematical properties, can be particularly useful in understanding the extended structure. It should be noted however that a detailed model of the intervening lens must be available to perform the transfer from the image plane (observational data) to the source plane (intrinsic properties of the quasar). We applied this technique for the first time to recover the properties of the molecular torus in the Cloverleaf, a quasar at $z=2.56$ (Kneib et al.\\ 1998), comparing HST images and IRAM interferometer CO maps (Alloin et al.\\ 1997). The CO-emitting region in the quasar was found to be a disk or ring-like structure orbiting the central engine at a radius between 75 and 100 pc, with Keplerian velocity around 100 \\kms . The effective resolution resulting from this technique turned out in this case to be about 20 times smaller than the synthesized beam size of the CO interferometer data.    Such significant benefits -- flux boosting and increased effective angular resolution -- have led us to focus our attention on gravitationally lensed systems and to conduct a CO survey of these objects. Another possible benefit of lensed versus unlensed objects is the potential for higher flux boosting for higher-$J$ transitions (differential magnification).  Since experience has shown that the best selection criterion for CO detection is the presence of detectable far-IR or submm/mm dust continuum emission, we started with a continuum survey of the known lensed quasars using the IRAM 30m radio telescope and the JCMT (Barvainis \\& Ivison 2002). A high dust continuum detection rate encouraged us to pursue a CO search with the IRAM interferometer (Guilloteau et al.\\ 1992).  Since this project was started early in the continuum survey, we were not at that time able to make a general selection based upon submillimeter flux.  Instead, the sample consisted of most of the then-known lensed quasars having optical redshifts measured to good accuracy. We observed 18 gravitationally lensed quasars, with redshifts in the range $1.375-3.595$.  However, reliable systemic redshifts remain a major difficulty for CO searches at high $z$ because currently available redshifts are mostly derived from highly ionized species in the quasar broad line region. As this region is often coupled to a high velocity wind, redshifts derived this way have been found to be blueshifted up to 1200~\\kms with respect to the systemic velocity of the host galaxy and the molecular environment of the quasar probed by CO measurements.  A typical offset is  600 \\kms , but there is wide dispersion from one object to another. Meanwhile, spectrometer bandwidths in the millimeter domain are too narrow ($\\sim 1500$~\\kms\\ at 3mm) to fully span this redshift uncertainty using a single central frequency setting. The combination of these two facts makes it likely that some CO lines will be missed in the course of a survey. In the case of the present survey, whenever the quasar redshift was from highly ionized species we applied a 600 \\kms redshift increment to search for its CO emission. We are fully aware that this offset, although statistically meaningful, may be just incorrect for some individual quasars.  In Sect.\\ 2 we describe the sample of gravitationally lensed quasars and the acquisition and reduction of the interferometer data set. Results, both in CO line emission and in the 1mm and 3mm continua, are also presented in Sect.\\ 2 for the entire sample. In Sect.\\ 3, we discuss the general results of the CO survey, and in Sect.\\ 4 consider the detection of MG0751+2716 in the CO(4$-$3) transition in more detail. Conclusions and future prospects are given in Sect.\\ 5. ",
        "conclusions": "Though the present survey yielded a low detection rate in CO, there are several new lensed quasar candidates yet to be observed based on their strong 850~$\\mu$m continua, recently discovered in the course of the submillimeter survey by Barvainis \\& Ivison (2002).  Supplementary, expanded-frequency observations of some sources (most notably RXJ0911+055) may turn up more CO detections from the present source list.  As for MG0751+2716, the centimeter radio source has four components connected by arcs (Leh\\'ar et al.\\ 1997), and in the optical it appears as a $1\\arcsec$ diameter Einstein Ring.  A primary driver for this project was to find lensed sources that could be spatially resolved in CO line emission.  This is currently possible for MG0751+2716 using the PdBI, and, like the Cloverleaf, reconstruction of the molecular source structure and kinematics on very small angular scales using the lensing properties may prove to be quite interesting."
    },
    "0201/astro-ph0201103_arXiv.txt": {
        "abstract": "We investigated ram-pressure effects by an intracluster wind on an inner disk of spiral galaxies by a hydrodynamical simulation. Even if the wind is mild and not strong enough to strip the gas disk, the ram pressure disturbs orbits of the inter-arm gas significantly. This results in asymmetric dense molecular arms in the inner few kpc region of a galaxy. This mechanism would explain the asymmetric CO gas distributions in the central regions often observed in Virgo spirals.  Key words: galaxies:   spiral --- galaxies: kinematics and dynamics --- galaxies: ISM --- galaxies: cluster of --- intergalactic matter ",
        "introduction": "Ram-pressure by intracluster medium (ICM) causes the stripping of interstellar matter (ISM) from galaxies (Farouki,  Shapiro 1980; Kritsuk 1984; Gaetz et al. 1987; Balsara et al. 1994; Sofue 1994; Quilis et al. 2000). It also produces a disturbed distribution of ISM in galaxies, such as head-tail H {\\sc i} structures as observed in Virgo galaxies (Cayatte et al.  1990; Vollmer et al. 2000, 2001; Phookun et al. 1993, 1995). The ram-pressure effect has, thus, been discussed mainly in relation to \\HI\\ gas stripping and outer disk structures. However, little attention has been paid to its effect on the inner molecular disk and arms. Only a few authors have discussed the ram effect on the spiral structure (Tosa 1994) and molecular clouds (Kenney et al. 1990; Sofue 1994).  Kenney et al. (1990) have shown that Virgo galaxies often exhibit asymmetric inner molecular disks. A recent high-resolution CO-line survey of Virgo cluster galaxies such as NGC 4254 and NGC 4654 (Sofue et al. in preparation) has revealed  that some of them show a significant asymmetry of the inner molecular disks and arms. It is not clear if such an inner deformation of dense gas disks can be produced by ram pressure effect, which is thought to be the cause of their deformed \\HI\\ envelops. Since these two galaxies have no massive companion that can disturb such inner disks, and they both show a head-tail \\HI\\ outer structure (Phookun et al. 1993, 1995), the inner deformation of molecular disks could also be due to the ram-pressure effect, while such an inner ram effect has not yetbeen investigated.  If the wind is very strong, such as assumed by Quilis et al. (2000) for the core of a rich cluster with an ICM density of $\\sim 3 \\times 10^{-3}$ atoms cm$^{-3}$ and  a wind velocity higher than $\\sim2000$ \\kms, the ISM of any galaxies would be completely stripped. On the other hand, if the wind is mild, in such a case as for the Virgo cluster, where the ICM density is of the order of $10^{-3} - 10^{-4}$ and the velocity is $\\sim 1000$ \\kms, outer \\HI\\ envelopes are deformed to produce head-tail structures (Vollmer et al. 2000, 2001). In current simulations, such as those by Vollmer et al., which were aimed at gas stripping and tailing of the outer \\HI\\ disks, the detailed behavior of the inner disk gas inside $\\sim 10$ kpc was not well understood because of the resolution.  In the present paper, we consider a mild ICM wind, and discuss its effect on the inner disk gas based on 2D hydrodynamical simulations with higher resolution than those aimed at outer \\HI\\ stripping, as above. If we simply apply the ram-stripping condition to an azimuthally structure-less gas disk, the ram pressure would hardly affect the inner disk. However, if we consider a spiral structure with an arm-to-interarm density contrast, it may happen that the ram-pressure can affect the low-density interarm gas. We consider here the possibility that the ram-pressure can affect the dense molecular gas within the central  few kpc region of spiral galaxies through disturbances of the orbits of inter-arm gas, even if the ICM wind is not strong enough to strip the disk gas. ",
        "conclusions": ""
    },
    "0201/astro-ph0201273_arXiv.txt": {
        "abstract": "\\baselineskip=11pt Many adaptive optics systems operate by measuring the distortion of the wavefront in one wavelength range and performing the scientific observations in a second, different wavelength range. One common technique is to measure wavefront distortions at wavelengths $<\\sim$1 $\\mu$m while operating the science instrument at wavelengths $>\\sim$ 1 $\\mu$m.  The index of refraction of air decreases sharply from shorter visible wavelengths to near-infrared wavelengths. Therefore, because the adaptive optics system is measuring the wavefront distortion in one wavelength range and the science observations are performed at a different wavelength range, residual image motion occurs and the maximum exposure time before smearing of the image can be significantly limited.    We demonstrate the importance of atmospheric differential refraction, present calculations to predict the effect of atmospheric differential refraction, and finally discuss the implications of atmospheric differential refraction for several current and proposed observatories. ",
        "introduction": "Adaptive Optics (AO) has been used for astronomical observations for more than a decade and numerous medium-to-large telescopes around the world are now equipped with AO systems.  If an AO system's wavefront sensor operates at different wavelengths than are being observed by the science instrument, and no correction is made for atmospheric differential refraction, the target object will appear to drift with respect to the science instrument.  An example of the problem this phenomenon can introduce is shown in Fig.~\\ref{fig:titanimage} which is the difference of two images of Saturn's moon Titan taken just 2.5 minutes apart while continuously tracking and correcting on Titan with the Keck II telescope's AO system. The maximum exposure time possible without degrading the spatial resolution of the data is significantly restricted if no correction for atmospheric differential refraction is made.   Adaptive optics systems for astronomical observing operate by splitting the incoming light into two beams, one of which goes to the wavefront sensor of the AO system and the goes to the science instrument.  In many cases a dichroic optic is used as the beamsplitter in order to send visible light ($<\\sim1 \\mu$m) to the wavefront sensor and infrared light ($>\\sim1 \\mu$m) to the science instrument. Examples of telescopes with AO systems that can operate in this way include the 3-m Shane telescope at Lick Observatory \\citep{2000SPIE.4007...63G}, the CFHT 3.6-m telescope \\citep{1998PASP..110..152R}, the Gemini \\begin{figure}[h] \\vskip 0in \\centerline{\\epsfig{figure=fig1.eps,width=2.6in,angle=0}} \\vskip 0in \\caption{\\small \\baselineskip=10pt The difference between two raw images of Titan taken just 153 seconds apart with NIRSPEC's SCAM detector on the Keck II telescope while continuously guiding/correcting with the AO system on Titan. Each image is the result of three 10 second exposures.  At the start of the first sequence of exposures Titan was at an elevation of 57$\\fdg$86 and an hour angle of 2.25; 153 seconds later when the second set of exposures started Titan was at an elevation of 57$\\fdg$26.  Titan's declination was 16$\\fdg$73. The 1-pixel wide slit of NIRSPEC's spectrometer is across Titan's disk and the apparent motion of Titan due to the effect of atmospheric differential refraction is obvious along the elevation vector. \\label{fig:titanimage}} \\end{figure} North 8-m telescope \\citep{2000SPIE.4007...26G}, and both of the W.M.~Keck Observatory's 10-m telescopes \\citep{2000PASP..112..315W}. The wavefront sensor measures the distortions to the wavefront and a correction is calculated and applied to a deformable mirror. The first order of distortion that an AO system is called upon to correct is simply image motion, commonly known as `tip/tilt'.  If significant tip/tilt residuals remain after AO correction then the resulting data will be of lower spatial resolution, regardless of how well the AO system corrects for the higher order terms of focus, astigmatism, coma, and so forth.  Minimizing tip/tilt residuals is critical to achieving optimum performance from an AO system.  The effect of atmospheric differential refraction is to introduce a systematic tip/tilt error that an AO system will not correct for unless specifically accounted for in the AO control software.  For ease in the remainder of this report we will refer to atmospheric differential refraction as ``ADR''.  Problems due to ADR arise when the AO system is correcting on visible wavelength light and the science instrument is observing at infrared wavelengths. From visible to infrared wavelengths the refractive index of air decreases sharply, as shown in Fig.~\\ref{fig:indexofrefraction}. A star's visible pointing center always appears at higher elevation than a star's infrared pointing center, except when the star is at the zenith.  That the two pointing centers do not coincide is not unto itself a problem for the typical AO system, however the offset between the two pointing centers is not constant. Without some consideration for the effect of ADR, a properly performing AO system will hold the visible pointing center of the star fixed, relative to both the wavefront sensor and the science camera.  However as time progresses and the star moves in elevation, the infrared pointing center will drift with respect to the visible pointing center, and thus in the infrared the star appears to drift with respect to the science instrument. Observers must consider this effect in determining the maximum exposure times in order to avoid `trailed' images, unless some other compensation is made.  \\begin{figure}[h] \\vskip 0in \\hskip 0.3in \\epsfig{figure=fig2.eps,width=2.4in,angle=0} \\vskip 0in \\caption{\\small \\baselineskip=10pt The refractive index of air as a function of wavelength across the visible and near-infrared spectrum at standard temperature and pressure in the absence of water vapor. \\label{fig:indexofrefraction}} \\end{figure}  In this paper we present calculations demonstrating when and how much of a problem ADR can be for an AO system.  Using data taken with the AO system on the Keck II telescope we show that in a long exposure, or sequence of exposures, the most significant uncorrected tip-tilt motion is due to the effect of ADR.  We further show that the effect of ADR is dependent on the spectral type, or color, of the star being used as a reference source. Finally, we discuss the implications of ADR for AO observing with current and currently proposed telescopes.  The code used in these calculations is available with the electronic version of this paper or by request from the author. While the code is written in the commonly used IDL programming language, it could easily be translated to other languages.  We encourage others to investigate the implications of ADR for their favourite telescope, site, target, observing strategy, etc.. ",
        "conclusions": "Using theoretical calculations and data from the W.M.\\ Keck II telescope we demonstrate that atmospheric differential refraction (ADR) should be considered when designing and building adaptive optics (AO) systems.  We present calculations and IDL code for others to calculate the effect of ADR for their own particular observing parameters.  The primary effect of ADR on typical AO observations is to reduce the maximum exposure time possible without significant image blurring.  Maximum exposure time decreases approximately linearly with increasing telescope size. Due to the variation of the refractive index of air across visible wavelengths, the maximum exposure time is strongly a function of the effective wavelength of the wavefront sensor.  The other important parameters in calculating maximum exposure time are observatory altitude and latitude, target declination and hour angle, and, to a lesser extent during typical near-infrared science observations, the effective wavelength of the science observations.  Planning for AO systems on larger ($>10$-meter) future telescopes must include consideration for how to compensate for the effect of ADR."
    },
    "0201/astro-ph0201045_arXiv.txt": {
        "abstract": "We investigate properties of astrometric microlensing of distant sources (such as QSOs and radio galaxies) caused by stars in the Galaxy, mainly focusing on application to the VERA (VLBI Exploration of Radio Astrometry) project. Assuming typical parameters for the Galaxy disk and bulge, we show that the maximum optical depth for astrometric shift of 10 $\\mu$as-level is $8.9\\times 10^{-2}$ for QSO-disk lensing case and $3.8\\times 10^{-2}$ for QSO-bulge lensing case. We also find that the maximum optical depth for QSO-disk lensing is larger by an order of magnitude than that for disk-disk or bulge-disk lensing case (assuming a typical source distance of 8 to 10 kpc). In addition to optical depth, we also calculate the event rate and find that the maximum event rate for QSO-disk lensing case is 1.2$\\times 10^{-2}$ event per year, which is about 30 times greater than that for disk-disk lensing case. This high event rate implies that if one monitors 10 QSOs behind the Galactic center region for 10 years, at least one astrometric microlensing event should be detected. Therefore, monitoring distant radio sources with VERA can be a new tool to study astrometric microlensing caused by stars in the Galaxy.  We also study the event duration of astrometric microlensing, and find that the mean event duration for QSO-disk lensing is 7.5 yr for QSOs located near the galactic center. This event duration for QSO-disk lensing is reasonably short compared to the project lifetime of VERA, which is anticipated to be $\\sim$ 20 yr. We also find that while the minimum event duration for bulge-bulge lensing is as short as 2.6 yr, the event duration for disk-disk lensing case cannot be shorter than 15 yr. Thus, although astrometric microlensing of bulge sources/lenses can be studied by optical astrometric missions like SIMA and GAIA, detections of disk events with the space astrometric missions are fairly difficult because of the limited project lifetime (typically $\\sim$5 yr) as well as the heavy dust extinction. Therefore, for studying astrometric microlensing by disk stars, VERA can be one of powerful tools by observing distant sources like QSOs and radio galaxies. We discuss the implications of astrometric microlensing for VERA by focusing on estimating the lens mass, and also present some possible candidates of radio sources toward which astrometric microlensing events should be searched for with VERA. ",
        "introduction": "Gravitational microlensing is one of the most promising tools to study invisible lenses like MACHOs and low-mass stars in the Galaxy. Paczynski (1986) first proposed to use microlensing effect to search for MACHOs in the Galaxy's halo based on a photometric monitoring of millions of stars in the Galaxy's bulge and the Magellanic Clouds. A number of groups have been conducting such `photometric' searches of microlensing to detect source magnification by gravitational microlensing (e.g., MACHO, EROS, OGLE, MOA, etc.), finding more than a hundred of photometric microlensing events.  In addition to such photometric microlensing, there is another type of microlensing called `astrometric microlensing', in which the positional shift of the lensed image, rather than the magnification of the source, is used to detect microlensing events. Recently, a number of studies have been made on such `astrometric' microlensing events (e.g., Hosokawa et al. 1993; 1997; H\\o{}g, Novikov \\& Polnarev 1995; Miyamoto \\& Yoshii 1995; Walker 1995; Paczynski 1996; 1998; Miralda-Escude 1996; Boden et al. 1998; Dominik \\& Sahu 2000). One of the major findings of these studies is that the probability of astrometric microlensing is much larger than that of photometric microlensing (e.g., Miralda-Escude 1996; Hosokawa et al.1997; Dominik \\& Sahu 2000; Honma 2001). For instance, the size of a photometric lens is equal to the Einstein-ring size given by \\begin{equation} \\label{eq:E-ring} R_{\\rm E}=\\sqrt{\\frac{4GM}{c^2} \\frac{D_{\\rm d}D_{\\rm ds}}{D_{\\rm s}}}. \\end{equation} Here $M$ is the lens mass, and $D_{\\rm d}$, $D_{\\rm ds}$, and $D_{\\rm s}$ are the observer-lens distance, lens-source distance, and observer-source distance, respectively. If a source comes within this radius from the lens, a source is magnified by more than a factor of 1.34. Meanwhile, the size of an astrometric lens is larger than $R_{\\rm E}$ by a factor of $\\beta_{\\rm max}$, which is given by following equation (e.g., Honma 2001). \\begin{equation} \\label{eq:beta_max} \\beta_{\\rm max} = \\theta_{\\rm E}/\\theta_{\\rm min}. \\end{equation} Here $\\theta_{\\rm E}$ is the angular size of the Einstein-ring radius ($\\theta_{\\rm E}\\equiv R_{\\rm E}/D_{\\rm d}$), and $\\theta_{\\rm min}$ is the minimum angular shift that can be detected by astrometric observation. Thus, if a source comes within $\\beta_{\\rm max} R_{\\rm E}$ from the lens, the source position is shifted by larger than $\\theta_{\\rm min}$. For astrometric lensing, $\\beta_{\\rm max}$ can be as large as 50 $\\sim$ 100 assuming $\\theta_{\\rm min}$ of 10-$\\mu$as level (e.g., Honma 2001).  Such a high astrometric accuracy, although not yet achieved, will be available soon because a number of astrometric missions are planned in early 21st centuries. For instance, there are four space astrometric missions to be launched by $\\sim$2010: DIVA, FAME, SIM and GAIA. Astrometric accuracies anticipated for those missions are: 200 $\\mu$as for DIVA, 50 $\\mu$as for FAME, and 10 $\\mu$as (or possibly higher) for SIM and GAIA. In addition to those space missions, there is another ground-based astrometric mission called VERA (VLBI Exploration of Radio Astrometry, e.g., Sasao 1996; Honma, Kawaguchi \\& Sasao 2000 and reference therein), which utilize a phase-referencing VLBI technique for astrometry of radio sources. The most remarkable difference between VERA and space astrometric missions is that VERA can observe thousands of distant sources like QSOs and radio galaxies to trace effect of astrometric microlensing. An advantages of using distant radio sources is that the column density of lens can be much higher, leading to higher event probability.  Investigating disk stars by astrometric microlensing with missions described above has significant scientific merits. First, astrometric microlensing provide us a new tool to study the lower end of stellar mass function, which is thought to be dominant populations in the Galaxy's disk. To date, the shape of the mass function remains uncertain in particular at its lower end, because the lower main-sequence stars are too faint for optical observations. Studying the low-mass star may also inform us about the nature of dark matter in the disk, as the disk dark matter may be baryonic matter in the form of low-mass stars and/or brown dwarfs. Further, studies of astrometric microlensing by disk stars will also provide fundamental information on parameters of the Galaxy's disk. For instance, one can estimate the disk scale length from the optical depth distribution with the Galactic latitude $l$ (e.g., Dominik \\& Sahu 2000). Also, it is possible to extract the disk density from astrometric microlensing events provided sufficient number of events are observed for statistics. From the scale length and disk density, one can estimate the total mass of the Galaxy's disk, and thus the density profile of dark halo inside in the disk region can be also deduced.  While the astrometric microlensing of stars are extensively studied in previous studies focusing on implications for SIM and GAIA(e.g., Dominik \\& Sahu 2000), there has been no study on astrometric microlensing of distant radio sources by disk stars considering an application to VERA. The most remarkable aspect of using distant sources is that some important parameters such as the lens mass may be determined independently of lens distance (e.g., Honma 2001), as well as high event probability due to large column density toward a source. For these reasons, in the present paper we extensively study the astrometric microlensing of distant sources by disk stars, and discuss its implications for VERA.  The plan of this paper is as follows: in section 2 we present calculations of optical depth, and in section 3 event rate. We also compare results for QSO-disk/bulge lensing cases with disk/bulge-disk lensing cases. In section 4, we present event durations of astrometric events, and in section 5 we discuss how the lens mass can be estimated from event durations. Finally in section 6, we discuss the implication of astrometric microlensing for VERA, and also summarize the findings of the present paper. ",
        "conclusions": "In this section, we discuss the implication of astrometric microlensing for VERA by focusing on the strategy of practical observations. First, here we briefly summarize advantages as well as disadvantages of using VERA for astrometric microlensing search. Advantages of the distant source case (which is the case for VERA) can be summarized as follows: I) large event probability expected from the optical depth and event rate, II) short event duration compared to the mission lifetime, and III) less uncertainty in lens mass determination, as the expected event duration is independent of the lens distance. On the other hand, disadvantages of distant source observation with VERA are: I) small number of sources compared to stars in the Galaxy, and II) possible structural variation of the sources which could cause an apparent position shift of sources.  As for the first disadvantage (shortage of sources), it is true that the number of radio sources is fairly small compared to the number of visible stars in the Galaxy. For instance, the number of compact radio sources that are expected to be observable with VERA is around 2000 (e.g., Ma et al. 1998, Peck \\& Beasley 1998). Moreover, the number density of compact sources in the Galactic plane is smaller than that in the off-plane region, since surveys of compact sources are conducted mainly in the off-plane regions (e.g., Patnaik et al. 1992; Peck \\& Beasley 1998). Thus, in order to have a sufficient number of sources, one has to conduct a survey in the Galactic plane region (e.g., Honma et al. 2000).  However, even at this stage there exist some sources which can be used for astrometric microlensing search. For instance, there is a radio source only 0.7$^\\circ$ away from the Galactic center (Backer \\& Sramek 1999; Reid et al. 1999). Also, according to Lazio \\& Cordes (1998) there are at least five extra-galactic radio sources within a few degrees from the Galactic center (including the one 0.7$^\\circ$ away from the Galactic center). Properties of the five sources are listed in table 3, including event rates for disk lens and bulge lens cases. The event rates for disk lens cases vary from $6.0\\times 10^{-3}$ to $1.1\\times 10^{-2}$ event/yr, and the total event rate for the five sources is found to be $4.0\\times 10^{-2}$ event/yr. Hence, even if one monitors only these five sources, one would detect an astrometric microlensing event caused by a disk star within 25 yr. The event rate for bulge lens is also quite high, varying from 2.8 to 4.7 $\\times 10^{-3}$ event per year. If the bulge event rate is included, the total event rate for the five sources is $6.0\\times 10^{-2}$ event/yr, indicating that within 16 years at least one of the sources is being lensed by bulge or disk star. Note that these five sources are located in about $2^\\circ \\times 8^\\circ$ area around the Galactic center (see figure 2 in Lazio \\& Cordes 1998). If the same number density applies to the rest of the Galactic plane, one can expect to find fairly large number of radio sources (i.e., a few tens). Therefore, the shortage of sources can be resolved by further search for radio sources in the Galactic plane.  As to the second disadvantage (QSO structure variation), it is not easy to distinguish the structure effect and astrometric microlensing within a short period. However, the image trajectory of an astrometric microlensing is a perfect circle for most cases except for events with extremely small impact parameter. On the other hand, it is unlikely that an image motion caused by the structure variation becomes a perfect circle. Hence, if one monitors the source position for the full event duration, one can easily discriminate the structural effect and astrometric microlensing (but of course there is no way to separate both effects if the two effects simultaneously happen to one source, and to avoid this it is better not to observe sources with significant structure variation).  In conclusion, astrometric microlensing events due to stars are practically detectable based on the observation of distant radio sources with VERA. If such events are detected, one can estimate the lens mass, and this will probably brings us new information on the stellar mass function at the lower end. According to the event rate calculation presented in the previous sections, it is not easy to detect an event, but it is possible if one monitors a few tens of sources for a decade or so. Although this is time-consuming, it is a worthwhile study when one considers its scientific importance.  \\newpage"
    },
    "0201/astro-ph0201267_arXiv.txt": {
        "abstract": "{\\xmm observed the Seyfert 1.9 galaxy \\n4258 in December 2000. At energies above 2~keV a hard nuclear point source is resolved that can be fitted by a highly absorbed power-law spectrum (\\nh = (8.0$\\pm$0.4)\\hcm{22}, photon index 1.64$\\pm$0.08) with an unabsorbed luminosity of 7.5\\ergs{40} in the (2--10) keV band. No narrow iron K$\\alpha$ emission line is detected (90\\% upper limit of equivalent width EW $\\sim$40 eV). The nuclear emission flux was observed to remain constant over the observation. A short archival \\chandra observation taken in March 2000 further constrains the hard emission to a point source coincident with the radio nucleus. A point source $\\sim$3\\arcsec\\ southwest of the nucleus does not contribute significantly. Spectral results of the \\chandra nuclear source are comparable (within the limited statistics) to the \\xmm parameters. The comparison of our iron line upper limit with reported detections indicates variability of the line EW. These results can be explained by the relatively low nuclear absorption of \\n4258 (which is in the range expected for its intermediate Seyfert type) and some variability of the absorbing material. Reflection components as proposed to explain the large iron line EW of highly absorbed Seyfert 2 galaxies  and/or variations in the accretion disk are however imposed by the time variability of the iron line flux. ",
        "introduction": "\\n4258\\,(\\object{M106}) is a B$_{\\rm T}$ = 8.5 mag nearby \\citep[distance of 7.2 Mpc, i.e. 1$\"\\cor35$~pc,][]{1999Natur.400..539H}, highly inclined \\citep[72\\degr, ][]{1988ngc..book.....T} SABbc spiral spectroscopically classified as a 1.9 Seyfert galaxy \\citep{1997ApJS..112..315H}.  The strong polarization of the relatively broad optical emission lines \\citep{1995ApJ...455L..13W} further supports the existence of an obscured active nucleus in \\n4258. Water maser line emission of rotating gas near the center of \\n4258 is modeled by a 3.6\\expo{7} M$_{\\sun}$ mass within 0.13 pc of the nucleus, indicating a massive nuclear black hole inside a highly inclined (83\\degr) thin gaseous disk \\citep{1995Natur.373..127M,1995ApJ...440..619G,1995A&A...304...21G}.  \\citet{1994ApJ...421..122T} detected a very bright radio continuum nuclear source, showing, in observations from March 1985 through to May 1990, long term time variability by 47\\% both at 6 and 20 cm \\citep{2001ApJ...551..702H}.  In X-rays ROSAT PSPC and HRI observations of \\n4258 \\citep{1994A&A...284..386P,1995ApJ...440..181C,1999A&A...352...64V} resolved diffuse emission mostly connected to the ``anomalous'' arms and emission from \\n4258 point sources. However, due to the soft energy band (0.1--2.4 keV) only upper limits could be derived for emission from the nucleus.  ASCA observations of \\n4258 were of reduced spatial resolution but extended the energy coverage to 10 keV with improved spectral resolution. \\citet{1994PASJ...46L..77M} reported, as well as soft emission components, a hard, highly absorbed point-like component that could be modeled by an absorbed power-law (\\nh\\ $\\sim1.5$\\hcm{23}, photon index $\\sim$1.78, unabsorbed luminosity 4\\ergs{40} in the (2--10) keV band) and an iron K emission line with an equivalent width of (250$\\pm$100) eV. Further ASCA observations improved on the spectral parameters of the hard component, confirming a narrow iron K$\\alpha$ emission line at ($6.45^{+0.10}_{-0.07}$) keV with an equivalent width of ($107^{+42}_{-37}$) eV, and showing time variability of the (5--10) keV flux and probably also the absorbing column of the hard component \\citep{2000ApJ...540..143R}.  \\sax observations extended the spectral coverage of the power-law component to 70 keV \\citep{2001ApJ...556..150F}. \\citet{2001ApJ...560..1W} have recently reported the results of a \\chandra observation optimized to investigating the X-ray emission from the ``anomalous arms\". The region around the nucleus showed two compact sources, a bright heavily absorbed source coincident with the radio nucleus and the nuclear H$_2$O maser source, and a second fainter source offset to the SW by 2\\farcs5. While the nuclear source is heavily piled-up, debilitating the determination of reliable spectral parameters, the spectrum of the weaker source is well described by an absorbed power-law (\\nh\\ $\\sim2$\\hcm{21}, photon index $\\sim$1.5, unabsorbed luminosity 5\\ergs{38} in the (0.5--4.5) keV band), and is most likely an X-ray binary within \\n4258 itself.  An \\xmm observation of \\n4258 was carried out to investigate the different emission components.  In this paper we will concentrate on the hard X-ray emission from the active nucleus and defer an analysis of the \\n4258 diffuse emission and extra-nuclear point sources to a later paper.  We also analyzed a short \\chandra ACIS-S observation from the \\chandra archive adding to the nuclear point source interpretation of the hard emission. ",
        "conclusions": "\\begin{table*} \\caption[]{Historical high energy spectral data for \\n4258.} \\begin{tabular}{lrrrrrccr} \\hline\\noalign{\\smallskip} \\multicolumn{1}{l}{Observatory}  & \\multicolumn{1}{c}{Date} & \\multicolumn{1}{c}{\\nh} & \\multicolumn{1}{c}{$\\Gamma$} & \\multicolumn{1}{c}{$E_{\\rm line}$} & \\multicolumn{1}{c}{$W_{\\rm line}$} & \\multicolumn{1}{c}{$f_{\\rm X}^{ *}$}    & \\multicolumn{1}{c}{$L_{\\rm X}^{ **}$}    & \\multicolumn{1}{c}{Ref.}\\\\ \\noalign{\\smallskip} & & \\multicolumn{1}{c}{[\\ohcm{22}]} &   & \\multicolumn{1}{c}{[keV]} & \\multicolumn{1}{c}{[eV]} &  &   & \\\\ \\noalign{\\smallskip}\\hline\\noalign{\\smallskip} ASCA     & May 1993 & $15\\pm2$             & $1.78\\pm0.29$          & $6.5\\pm0.2$ & $250\\pm100$ &     & 4.2  &  1   \\\\ ASCA     & May 1993 & $13.6^{+2.1}_{-2.2}$ & $1.78^{+0.22}_{-0.26}$ &             &             & 5.1 &      &  2   \\\\ ASCA     & May 1996 & $9.2\\pm0.9$          & $1.71^{+0.18}_{-0.17}$ &             &             & 8.3 &      &  2   \\\\ ASCA     & Jun 1996 & $8.8^{+0.7}_{-0.6}$  & $1.83\\pm0.13$          &             &             & 8.8 &      &  2   \\\\ ASCA     & Dec 1996 & $9.7\\pm0.8$          & $1.87\\pm0.15$          &             &             & 9.5 &      &  2   \\\\ \\sax     & Dec 1998 & $9.4\\pm1.2$          & $2.11\\pm0.14$          & $6.57\\pm0.20$ & $85\\pm65$ &     & 10   &  3   \\\\ ASCA     & May 1999 & $9.5^{+2.1}_{-0.9}$  & $1.86^{+0.40}_{-0.13}$ &             &             & 4.0 &      &  2   \\\\ ASCA     & May 1999 & $8.9^{+0.4}_{-0.7}$  & $1.83^{+0.06}_{-0.09}$ & $6.45^{+0.10}_{-0.07}$ &  $107^{+42}_{-37}$ &  & 5.7 &  2   \\\\ \\chandra & Mar 2000 & $7.2\\pm1.8$          & $1.4\\pm0.5$            &             &             &     & 13   & this work \\\\ \\xmm     & Dec 2000 & $8.0\\pm0.4$          & $1.64\\pm0.08$          &  6.45       & $<40$       &     &  7.5 & this work \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\end{tabular} \\label{spectra}  Notes and references:\\\\ $^{ *~}$: 5--10 keV uncorrected flux [\\oergcm{-12}]\\\\ $^{ **}$: 2--10 keV absorption corrected luminosity for \\n4258 distance of 7.2 Mpc [\\oergs{40}] \\\\ (1) \\citet{1994PASJ...46L..77M}; (2) \\citet{2000ApJ...540..143R}; (3) \\citet{2001ApJ...556..150F} \\end{table*}  ASCA observations of \\n4258 established a heavily absorbed hard component in the overall X-ray spectrum which was attributed to the nucleus and confirmed by \\sax \\citep{1994PASJ...46L..77M, 2001ApJ...556..150F}. The ASCA and \\sax observations of the hard emission of \\n4258 however, suffered from instrument PSFs of more than an arcminute, and only higher resolution instruments such as \\xmm and \\chandra were thought able to confirm a nuclear origin to the hard emission.  Even the \\xmm observations do not resolve the nuclear emission from that of another reasonably bright nearby point source, most likely an X-ray binary, that is resolved by \\chandra (see Sect. 3).  \\citet{2001ApJ...560..1W}, analysing a deeper \\chandra observation, report this source to be at a projected distance from the nucleus of 2\\farcs5 (87 pc) and characterize its spectrum as an absorbed power-law ($N_{\\rm H} = (2.0^{+1.2}_{-1.1})$\\hcm{21}, $\\Gamma = 1.49^{+0.50}_{-0.37}$, absorption corrected luminosity 5.1\\ergs{38}), similar in slope but much less absorbed than the nuclear spectrum.  Fortunately, the source is rather faint (less than 1\\% in flux) compared to the nuclear emission and cannot significantly influence the \\xmm results. Also we detect no other hard bright point-like sources within an arcmin of the nucleus that could significantly contribute to the ASCA and \\sax emission, which therefore really has to originate from the nucleus.  In Table \\ref{spectra} we give an historical record of the \\n4258 nuclear spectrum. The \\xmm power-law spectral parameters are within the range of the other measurements. The photon index is on the hard side but compatible within the errors of the ASCA results. Only the \\sax index pointed at a somewhat softer spectrum. The absorbing column ($\\sim$8\\hcm{22}) seems to have not varied significantly since the first ASCA measurements that pointed at an absorption a factor 1.5--2 higher. The luminosity however, seems to have varied on a timescale of several years by a factor of about two starting from low values in 1993 with a maximum in 1996 to 1998. The \\xmm measurements indicate that in December 2000 the nuclear luminosity was again at a lower level. \\citet{2001ApJ...556..150F} report variability of a factor of about 2 on half-day timescales as well as smaller scale variations (10\\%--20\\%) on timescales as short as 1 hr. During the \\xmm observations the source did not show any variability at similar amplitudes. With its absorption-corrected (2--10) keV luminosity of typically 7\\ergs{40}, the AGN in \\n4258 is presently not very active compared to other Seyfert galaxies. \\citet{1999ApJ...522..157R} report luminosities from  Seyfert galaxies ranging from  3\\ergs{40} to 2\\ergs{44}.  The high X-ray absorption of the nuclear component in \\n4258 is expected, given the optically-derived Seyfert 1.9 classification of the galaxy, using the so-called ``unified model'' \\citep{1993ARA&A..31..473A}. This model assumes that the engine is at work in all active galactic nuclei and differences between type 1 and type 2 AGN are ascribed solely to orientation effects, i.e. that our line of sight is (type 2) or is not (type 1) obscured by optically thick material.  \\citet{1999ApJ...522..157R} showed that most ``strict'' Seyfert 2 nuclei are highly obscured (\\nh $>$\\ohcm{23}) or even Compton thick(\\nh $>$\\ohcm{24}). Intermediate type 1.8--1.9 Seyfert galaxies on the other hand are characterized by an average \\nh\\ much lower than ``strict'' Seyfert 2 galaxies. The absorption of the \\n4258 nuclear spectrum of $\\sim$8\\hcm{22} fits into this scenario.  From the first reports of hard X-ray emission from \\n4258, a narrow iron K$\\alpha$ emission line was always needed to model the spectra. According to Table~\\ref{spectra}, the strongest line compared to the continuum was needed for the 1993 ASCA data. With \\xmm we do not detect a line, and our EW upper limit is more than a factor of 6 below the \\citet{1994PASJ...46L..77M} initial ASCA detection and still a factor of more than 2 below the later ASCA and \\sax values \\citep{2000ApJ...540..143R, 2001ApJ...556..150F}. Note that \\citet{1994PASJ...46L..77M} report that no background was subtracted from the May 1993 ASCA raw spectra and therefore, one might have doubts regarding the line strength in this analysis if the background contains contributions from an intrinsic iron feature. It must be stated however, that the existence of the line was clearly demonstrated in the later ASCA observations. The \\sax equivalent width from December 1998 is, due to the large errors, compatible with both the ASCA results from May 1999 and with our \\xmm upper limit from December 2000.  There are at least two possible explanations for this result: (1) The difference could be caused by the larger sky area used to accumulate the ASCA and \\sax spectra. In this way an extended component from the \\n4258 disk could be picked up that would not be covered by the much smaller spectral extraction region used here with \\xmm. We tried to rule out this explanation by extracting \\xmm spectra from a comparably large area. In these spectra we did not detect any iron line emission. The upper limit to the equivalent width however is similar to the measurements by earlier satellites. (2) The most straightforward explanation however, is that the line intensity is varying with time.  Several groups have calculated the contribution of an obscuring torus to the X-ray spectra of Seyfert galaxies using Monte Carlo simulations \\citep[e.g.][]{1991PASJ...43..195A, 1993MNRAS.263..314L, 1994MNRAS.267..743G}. They find strong dependence of the iron line EW at 6.4 keV on the absorbing column and geometrical factors. For the relatively low nuclear absorption in \\n4258, EWs of $\\sim$100 eV or less are expected. This is compatible with our findings. These simple models therefore would be sufficient to explain the \\xmm data and we would not need to invoke additional reflection components. Such a component is for instance needed to explain the iron line EW of $>$2 keV measured by \\chandra in the Seyfert 2 nucleus of M\\,51 \\citep{2001ApJ...560..139T}.  According to Table \\ref{spectra}, the iron line varied in absolute flux by a factor of at least 2 in 18 months (from May 1999 to December 2000) setting an upper limit on the size-scale of the fluorescing region of 0.5 pc. This seems to argue against a torus origin for the iron line which would imply a significantly larger size scale. Instead, it favours an origin of the line in the accretion disk much closer to the nucleus. If the line originated in the outer regions of the accretion disk as proposed by \\citet{2000ApJ...540..143R} using different arguments, its variability could reflect intensity changes of the nuclear power law source (that may not have been covered directly by the observations listed in Table \\ref{spectra}). Alternative explanations for the line variability might be changes in the disk (e.g. the ``inner'' edge of the optically thin disk, or the ionization structure of the disk surface).  The - up till now - sparse data on the \\n4258 6.4 keV line seem to indicate a correlation between the amount of absorption and the EW of the line. However, an anticorrelation of the line EW with the nuclear luminosity can also not be ruled out. More high quality spectra are clearly needed to clarify the behaviour, and such results may then force revisions and improvements to the simple unification picture."
    },
    "0201/astro-ph0201321_arXiv.txt": {
        "abstract": " ",
        "introduction": "Gamma-ray bursts (GRBs) are extremely bright ($10^{50}-10^{53}$~erg/s) and short ($10^{-2}-10^{2}$~s) emission events observed from distant parts of the Universe. Their redshifts are now measured in about 20 cases and typical $z\\sim 1$ (up to 4.5) are found~\\cite{Djorg}. The burst were observed to occur with a rate of $\\approx 1$ per day by the BATSE experiment~\\cite{Fishman} and even a higher rate would probably be detected with more sensitive instruments~\\cite{Stern1}. GRBs are very different from supernovae not only because of their short duration and high luminosity. Their most special feature is that the emission peaks in the gamma-ray band, at $h\\nu$ of a few hundred keV or perhaps more. In many cases ($\\approx 50$\\%) the gamma-ray bursts are followed by afterglows --- a much longer and softer emission that deacays in time and evolves from X-rays (hours) to radio (months)~\\cite{vanPar}.  What triggers the bursts is still uncertain. The primary energy release occurs in a very compact region (probably within $10^{7}$~cm --- from variability arguments) which is comparable to the size of compact objects --- black holes and neutron stars. The energy output (assuming isotropic explosion) approaches a stellar rest-mass energy. This suggests a high efficiency of mass conversion into radiation and points to a relativistic collapse with a gravitational potential $\\phi\\sim c^2$ which agrees with the potential of compact objects. Yet the configuration of the progenitor system and the reason of the collapse are uncertain (see reviews \\cite{Piran,Mesz1} and refs. therein). It can be (I) coalescence of a close binary consisting of two compact objects~\\cite{Pacz1}, (II) collapse of a massive star core~\\cite{Woosley}, and (III) collapse of a white dwarf~\\cite{Usov92}. Also conversion of neutron stars into strange stars was proposed~\\cite{Cheng,Bombaci}. Different scenarios can be tested against observations. In particular, massive stars spend their short lifes close to where they were born, and hence in the second scenario GRBs should occur in regions of active star formation. There is a growing evidence that this is indeed the case~\\cite{Djorg}.  \\begin{figure}[htb] \\begin{center} \\epsfig{file=dur.ps,height=9cm} \\caption{The duration distribution of 2041 GRBs from the current BATSE catalogue. Right panel: the duration is measured by T90, which is the time over which a burst emits from 5\\% of its total counts to 95\\%. Left panel: the duration is measured by $T_{50}$ --- the time over which a burst emits from 25\\% of its total counts to 75\\%. The counts are summed over all 4 BATSE energy channels (photon energy $E > 20$keV).} \\label{fig:dur} \\end{center} \\end{figure}  The afterglow is believed to come from a relativistic blast wave driven by the central explosion into an ambient medium. If GRBs are triggered by old stellar remnants they are likely surrounded by an interstellar medium (ISM) of density $n\\sim 1$~cm$^{-3}$ or even lower if the remnants are kicked out of the disk of the parent galaxy. In contrast, a GRB with a massive progenitor should occur inside a dense and strongly inhomogeneous star-forming cloud. Even more importantly, a massive progenitor emits a powerful wind all the time before it explodes, and this wind is the actual ambient medium of the GRB, with density scaling with radius as $R^{-2}$ out to a distance of a few parsecs~\\cite{Li}. The afterglow observations were supposed to reveal the nature of the ambient medium and thereby the progenitor. Yet after having observed many afterglows no definite conclusions are made. Perhaps the most surprising finding is that the afterglows are very diverse. ",
        "conclusions": "Owing to the recent observational progress, the state of the GRB field has crucially transformed. In the beginning of 90s, $\\sim 10^2$ theories did not contradict the data and were considered as principally possible; to a large extent the choice of a plausible theory was a matter of taste. The present situation is just opposite. GRBs have shown a mysterious complicated phenomenology with a number of well formulated observational facts. A successful theory is expected to make specific predictions that agree with the data on a quantitative level. The understanding of GRBs is far from such an ideal state, and the existing theories are rather naive and lacking a predictive power. Where models can fit the data (e.g. afterglow light curves) it is partially due to a large number of free parameters which ensure a sufficient flexibility of the model, and it does not make one confident that the simplifying assumptions are correct.  The difficulty of the GRB theory is well compensated by the recent observations which put more and more constraints. The observational progress may become even more impressive with additional channels of information such as neutrino and gravitational radiation. New exciting observations of electromagnetic radiation are expected from future missions --- Swift and GLAST. It allows one to hope for a future theory that would clarify the physics of the explosion and answer the basic questions: (1) What is the progenitor and why does it collapse? (2) Where is the released gravitational energy channeled to before it starts to feed the outflow (energy of rotation, magnetic energy, heat?) (3) What is the composition of the ejecta ($e^\\pm$, $p$, $n$, magnetic field?) and why are they so highly relativistic?"
    },
    "0201/hep-ph0201087_arXiv.txt": {
        "abstract": "Detailed predictions for the D-N asymmetry for the Super-Kamiokande experiment, as well as for the {\\it Full Night} and {\\it Core} D-N asymmetries in the solar neutrino induced CC event rate and the {\\it Day}, {\\it Night} and {\\it Core} ratios of the CC and NC event rates, measured in the SNO experiment, are derived in the cases of the LMA MSW and LOW solutions of the solar neutrino problem. The indicated observables for the SNO experiment are calculated for two values of the threshold (effective) kinetic energy of the final state electron in the CC reaction on deuterium: $T_{\\mathrm{e},\\mathrm{th}} = 6.75$\\ MeV and 5.0 MeV. The possibilities to further constrain the regions of the LMA MSW and LOW solutions of the solar neutrino problem by using the forthcoming SNO data on the D-N asymmetry and on the CC to NC event rate ratio are also discussed. ",
        "introduction": "\\vspace{-0.2cm}  \\hskip 0.5cm The recent SNO results \\cite{SNO1}, combined with the data from the Super-Kamiokande experiment \\cite{SKsol}, clearly demonstrate the presence of $\\nu_{\\mu}$ ( $\\nu_{\\tau}$) component in the flux of solar neutrinos reaching the Earth \\footnote{The non-electron neutrino component in the flux of solar neutrinos can also include, or correspond to, $\\bar{\\nu}_{\\mu}$ and/or $\\bar{\\nu}_{\\tau}$ \\cite{SNO1}.}. This represents a compelling evidence for oscillations and/or transitions of the solar neutrinos.  The SNO experiment measured the rate of the charged current (CC) reaction $\\nu_e + D \\rightarrow e^{-} + p + p$ for $T_e \\geq 6.75~{\\rm MeV}$, $T_e$ being the (effective) kinetic energy of the final state electron \\cite{SNO1}. The reaction is due to the flux of solar $\\nu_e$ from $^8$B decay having energy of $E \\gtap 8.2~{\\rm MeV}$. Assuming that the $^8$B neutrino energy spectrum is not substantially modified by the solar neutrino oscillations, the SNO collaboration obtained the following value of the solar $\\nu_e$ flux: \\begin{equation} \\Phi^{CC}(\\nu_e) = (1.75 \\pm 0.15) \\times 10^{6}~{\\rm cm^{-2}s^{-1}}, \\label{phinue} \\end{equation} \\noindent where we have added the statistical and systematic errors and the estimated theoretical uncertainty (due to the uncertainty in the CC reaction cross section) given in \\cite{SNO1} in quadrature. Utilizing the data on $\\Phi^{CC}(\\nu_e)$ and the data on the solar neutrino flux obtained by the Super-Kamiokande experiment, it is possible to deduce \\cite{SNO1} (see also \\cite{FogliSNO}) the value of the non-electron neutrino component in the flux of solar neutrinos measured by the Super-Kamiokande collaboration: \\begin{equation} \\Phi(\\nu_{\\mu,\\tau}) = (3.69 \\pm 1.13) \\times 10^{6}~{\\rm cm^{-2}s^{-1}}. \\label{phinumu} \\end{equation} \\noindent This flux is different from zero at more than 3 s. d.  Global analyses of the solar neutrino data \\cite{SKsol,Cl98,Kam96,SAGE,GALLEXGNO}, including the SNO results \\cite{SNO1} and the Super-Kamiokande data on the $e^{-}-$spectrum and day-night asymmetry, in terms of the neutrino oscillation hypothesis show \\cite{FogliSNO,ConchaSNO,GoswaSNO,StrumiaSNO,KSSNO,GiuntiSNO,MSmySNO1,ConchaSNO2,GiuntiSNO2,SKSmySNO} (see also \\cite{PeeSNO}) that the data favor the large mixing angle (LMA) MSW, the LOW and the quasi-vacuum oscillation (QVO) solutions of the solar neutrino problem with transitions into active neutrinos. In the case of the LMA solution, the range of values of the neutrino mass-squared difference $\\Delta m^2 > 0$, characterizing the two-neutrino transitions of the solar neutrinos into an active neutrino, $\\nu_e \\rightarrow \\nu_{\\mu(\\tau)}$, was found, e.g., in \\cite{FogliSNO} and \\cite{ConchaSNO} to extend (at 99\\% C.L.) to $\\sim 5.0\\times 10^{-4}~{\\rm eV^2}$ and $\\sim 8.0\\times 10^{-4}~{\\rm eV^2}$, respectively: \\begin{equation} {\\rm LMA~MSW}:~~~~~~2.0\\times 10^{-5}~{\\rm eV^2} \\ltap \\Delta m^2 \\ltap (5.0 - 8.0)\\times 10^{-4}~{\\rm eV^2}~. \\label{dmsolLMA} \\end{equation} \\noindent The best fit values of $\\Delta m^2$ obtained in the independent analyses \\cite{FogliSNO,ConchaSNO,GoswaSNO,StrumiaSNO,KSSNO} are grouped in the narrow interval $(\\Delta m^2)_{BFV} = (4.3 - 4.9)\\times 10^{-5}~{\\rm eV^2}$. A smaller best fit value was found in \\cite{ConchaSNO2}, $(\\Delta m^2)_{BFV} = (3.3 - 3.7)\\times 10^{-5}~{\\rm eV^2}$, while a larger value was obtained, e.g., in \\cite{SKSmySNO}: $(\\Delta m^2)_{BFV} = 6.0\\times 10^{-5}~{\\rm eV^2}$. Similar results, $(\\Delta m^2)_{BFV} = 6.3\\times 10^{-5}~{\\rm eV^2~and~} 6.1\\times 10^{-5}~{\\rm eV^2}$, were obtained in \\cite{GiuntiSNO} and in \\cite{GiuntiSNO2} by performing a Bayesian analysis of the solar neutrino data. For the mixing parameter $\\sin^22\\theta$, which controls the oscillations of the solar neutrinos, it was found, e.g., in \\cite{FogliSNO} at 99\\% C.L.: \\begin{equation} {\\rm LMA~~MSW}:~~~~~~~~~~~~~~~~~ 0.60 \\ltap \\sin^22\\theta \\ltap 0.99,~~~~~~~~~~~~~~~~~~~~~~~~~~~~~ \\label{thLMA} \\end{equation} \\noindent The best fit values of $\\sin^22\\theta$ obtained, e.g., in \\cite{FogliSNO,ConchaSNO,GoswaSNO,StrumiaSNO,KSSNO} are confined to the interval $(\\sin^22\\theta)_{BFV} = (0.79 - 0.82)$. Somewhat smaller values were found in \\cite{ConchaSNO2}, \\cite{GiuntiSNO2} and in \\cite{SKSmySNO}: $(\\sin^22\\theta)_{BFV} = (0.75 - 0.79);~0.76;~0.77$, respectively.  Detailed results were obtained in \\cite{FogliSNO,ConchaSNO,GoswaSNO,StrumiaSNO,KSSNO,GiuntiSNO,ConchaSNO2,GiuntiSNO2,SKSmySNO} for the LOW solution as well. The 95\\% C.L. allowed intervals of values of $\\Delta m^2$ and $\\sin^22\\theta$ found in \\cite{FogliSNO}, for instance, read: \\begin{equation} {\\rm LOW}:~~~~~~6.0\\times 10^{-8}~{\\rm eV^2} \\ltap \\Delta m^2 \\ltap 1.8\\times 10^{-7}~{\\rm eV^2}~,~~~ 0.94 \\ltap \\sin^22\\theta \\ltap 1.0. \\label{dmsolthLOW} \\end{equation} \\noindent The best fit values of $\\Delta m^2$  and $\\sin^22\\theta$ for the LOW solution, derived, e.g., in \\cite{FogliSNO,ConchaSNO,GoswaSNO,StrumiaSNO,KSSNO,ConchaSNO2} are compatible with each other and are all approximately given by $(\\Delta m^2)_{BFV} \\cong 10^{-7}~{\\rm eV^2}$ and $(\\sin^22\\theta)_{BFV} \\cong (0.94 - 0.97)$. A substantially different value of $(\\Delta m^2)_{BFV}$ was found in \\cite{SKSmySNO}: $(\\Delta m^2)_{BFV} \\cong 5.5\\times 10^{-8}~{\\rm eV^2}$ and $(\\sin^22\\theta)_{BFV} \\cong 0.99$.  The analyses \\cite{FogliSNO,ConchaSNO,GoswaSNO,StrumiaSNO,KSSNO,GiuntiSNO,SKSmySNO} were based, in particular, on the standard solar model (SSM) predictions of ref. \\cite{BPB01} (BP2000) for the different components of the solar neutrino flux ($pp$, $pep$, $^7$Be, $^8$B, $CNO$, $hep$, $^{17}$F). In \\cite{FogliSNO,ConchaSNO,GoswaSNO,StrumiaSNO,KSSNO,GiuntiSNO,ConchaSNO2} the published Super-Kamiokande data on the day-night (D - N) asymmetry \\cite{SKsol} were used as input in the analyses, while in \\cite{SKSmySNO} the latest (preliminary) results on the D-N asymmetry, obtained from the analysis of {\\it all currently available Super-Kamiokande solar neutrino data} was utilized (see further). The authors of ref. \\cite{ConchaSNO2} have used in their analysis a new value of the $^8$B neutrino flux which is suggested by the results of the latest (and more precise) experimental measurement \\cite{pBe701} of the cross section of the reaction $p + ^{7}$Be $\\rightarrow ^{8}$B $+ \\gamma$. According to the SSM, the $^8$B is produced in the Sun in the indicated reaction and the $\\beta^{+}-$decay of $^8$B in the central part of the Sun gives rise to the solar $^8$B neutrino flux. The results obtained in \\cite{pBe701} give a larger $p - ^7$Be reaction cross-section (with smaller uncertainty), and correspondingly - a larger astrophysical factor $S_{17}$ (see, e.g., \\cite{ConchaSNO2}) than the one used in \\cite{BPB01}, which implies, in particular, a larger value of the $^8$B neutrino flux than the value predicted \\footnote{The $^8$B neutrino flux predicted in \\cite{BPB01} reads $\\Phi(B)_{BP2000} = 5.05 \\times (1 ^{+0.20}_{-0.16})\\times 10^{6}~ {\\rm cm^{-2}s^{-1}}$, while the flux utilized in the analysis performed in \\cite{ConchaSNO2} is $\\Phi(B)_{NEW} = 5.93\\times (1 ^{+0.14}_{-0.13})\\times 10^{6}~ {\\rm cm^{-2}s^{-1}}$. } in \\cite{BPB01}. In the global Bayesian analysis performed in \\cite{GiuntiSNO2} the SSM predictions for the solar neutrino fluxes were not used: both the values of the fluxes and of the oscillation parameters were derived from the data.  The best fit values of $\\Delta m^2$ found in \\cite{FogliSNO,ConchaSNO,GoswaSNO,StrumiaSNO,KSSNO} differ from that derived in \\cite{SKSmySNO} essentially due to the difference in the Super-Kamiokande data on the D-N asymmetry used as input in the corresponding analyses: in \\cite{SKSmySNO} the latest (preliminary) Super-Kamiokande result implying a smaller mean value of the D-N asymmetry than the published one in \\cite{SKsol} was utilized. The smaller possible D-N asymmetry drives $(\\Delta m^2)_{BFV}$ to larger (smaller) value in the LMA MSW (LOW) solution region \\cite{SKSmySNO}. Although the data on the D-N asymmetry used in \\cite{FogliSNO,ConchaSNO,GoswaSNO,StrumiaSNO,KSSNO} and in \\cite{ConchaSNO2} are the same, the best fit value of $\\Delta m^2$ in the LMA MSW solution region found in \\cite{ConchaSNO2} is smaller than those found in \\cite{FogliSNO,ConchaSNO,GoswaSNO,StrumiaSNO,KSSNO} because of the difference between the values of the astrophysical factor $S_{17}$, and thus of the $^8$B neutrino flux, used in \\cite{ConchaSNO2} and in \\footnote{Let us note that, e.g., in \\cite{ConchaSNO,KSSNO,MSmySNO1,ConchaSNO2} results obtained by treating the $^8$B neutrino flux as a free parameter in the analysis were also reported. These results were taken into account when we quoted above the $\\Delta m^2$ and $\\sin^22\\theta$ best fit values. } \\cite{FogliSNO,ConchaSNO,GoswaSNO,StrumiaSNO,KSSNO}.   In the present article we update our earlier predictions \\cite{SK97I,SK97II,DNSNO00} for the D-N asymmetry for the Super-Kamiokande and SNO experiments, taking into account the recent progress in the studies of solar neutrinos. The day-night (D-N) effect - a difference between the solar neutrino event rates during the day and during the night, caused by the additional transitions of the solar neutrinos taking place at night while the neutrinos cross the Earth on the way to the detector (see, e.g., \\cite{HataL94,DNold} and the references quoted therein), is a unique testable prediction of the MSW solutions of the solar neutrino problem. The experimental observation of a non-zero D-N asymmetry \\begin{equation} A^{N}_{D-N} \\equiv \\frac{R_{N} - R_{D}}{(R_N + R_D)/2}, \\end{equation} \\noindent where $R_N$ and $R_D$ are, e.g., the one year averaged event rates in a given detector respectively during the night and the day, would be a very strong evidence in favor (if not a proof) of an MSW solution of the solar neutrino problem. Extensive predictions for the magnitude of the D-N effect for the Super-Kamiokande and SNO detectors have been obtained in \\cite{SK97I,SK97II,DNSNO00,LisiM97,BK97,SK98III,Barger01DN}. High precision calculations of the D-N asymmetry in the one year averaged recoil-e$^{-}$ spectrum measured in the Super-Kamiokande experiment and in the energy-integrated event rates for the two experiments were performed for three event samples, {\\it Night}, {\\it Core} and {\\it Mantle}, in \\cite{SK97I,SK97II,DNSNO00,SK98III}. The night fractions of these event samples are due to neutrinos which respectively cross the Earth along any trajectory, cross the Earth core, and cross only the Earth mantle (but not the core), on the way to the detector.  We focus here, in particular, on providing detailed predictions for the D-N asymmetry for the LMA MSW and the LOW solutions of the solar neutrino problem, which are favored by the current solar neutrino data. We will consider in what follows the {\\it Night} (or {\\it Full Night}) and the {\\it Core} D-N asymmetries, $A^{N}_{D-N}$ and $A^{C}_{D-N}$. The current Super-Kamiokande data \\cite{SKsol} do not contain evidence for a substantial D-N asymmetry: the latest published result on $A^{N}_{D-N}$ reads \\cite{SKsol} \\begin{equation} A^{N}_{D-N}(SK) = 0.033 \\pm 0.022~(stat.)~^{+0.013}_{-0.012}~(syst.), \\label{dnsk01} \\end{equation} \\noindent while the result of the latest analysis of {\\it all currently available Super-Kamiokande solar neutrino data} gives even smaller mean value \\cite{SKSmySNO} \\begin{equation} A^{N}_{D-N}(SK) = 0.021 \\pm 0.022~(stat.)~^{+0.013}_{-0.012}~(syst.). \\label{dnsk01Smy} \\end{equation} \\noindent Adding the errors in eqs. (\\ref{dnsk01}) and (\\ref{dnsk01Smy}) in quadrature, one finds that at 1.5 (2.0) s.d., $A^{N}_{D-N}(SK) < 0.072~(0.085)$ and $A^{N}_{D-N}(SK) < 0.060~(0.073)$, respectively.  We give in the present article also detailed predictions for another important observable - the ratio of the event rates of the CC reaction $\\nu_e + D \\rightarrow e^{-} + p + p$, $R_{SNO}(CC)$, and of the neutral current (NC) reaction $\\nu + D \\rightarrow \\nu + n + p$, $R_{SNO}(NC)$, induced by the solar neutrinos in SNO, \\begin{equation} R^{SNO}_{CC/NC} \\equiv \\frac{\\frac{R_{SNO}(CC)}{R_{SNO}(NC)}} {\\frac{R^{0}_{SNO}(CC)}{R^{0}_{SNO}(NC)}}~~, \\label{ccnc} \\end{equation} \\noindent which is normalized above to the value of the same ratio in the absence of oscillations of solar neutrinos, $R^{0}_{SNO}(CC)/R^{0}_{SNO}(NC)$. First results on the D-N asymmetry and on the CC to NC event rate ratio $R^{SNO}_{CC/NC}$ are expected to be published in the near future by the SNO collaboration. We discuss as well the possibilities to further constrain the regions of the LMA MSW and LOW-QVO solutions of the solar neutrino problem by using the forthcoming SNO data on the D-N asymmetry  $A^{N}_{D-N}$ and on the CC to NC event rate ratio $R^{SNO}_{CC/NC}$.  Updated predictions for the {\\it Night} D-N asymmetry and the {\\it average} CC to NC event rate ratio for the SNO experiment were derived after the publication of the first SNO results also in \\cite{KSSNO,ConchaSNO2}. However, our study overlaps little with those performed in \\cite{KSSNO,ConchaSNO2}. \\vspace{-0.5cm} ",
        "conclusions": "\\vskip -0.2cm \\hskip 0.5cm  In the present article we have derived detailed predictions for the D-N asymmetry in the solar neutrino induced CC event rate in the SNO detector for the LMA MSW and the LOW solutions of the solar neutrino problem, which are favored by the current solar neutrino data. We have obtained results for the {\\it Night} (or {\\it Full Night}) and the {\\it Core} D-N asymmetries for SNO, $A^{N}_{D-N}(SNO)$ and $A^{C}_{D-N}(SNO)$, which are presented in the form of iso-(D-N) asymmetry contour plots in the $\\Delta m^2 - \\tan^2\\theta$ plane in Figs. 2 - 3. Detailed predictions for the {\\it Night} and {\\it Core} D-N asymmetries for the Super-Kamiokande detector, $A^{N,C}_{D-N}(SK)$, were also derived (Fig. 1). The high precision calculations of $A^{N,C}_{D-N}(SNO)$ have been performed by taking into account, in particular, the energy resolution function of the SNO detector \\cite{SNO1}. Our results show, however, that the effect of the energy resolution function on the predicted values of the {\\it Full Night} and {\\it Core} D-N asymmetries is negligible when $A^{N,C}_{D-N}(SNO)\\geq 0.01$. The asymmetries $A^{N,C}_{D-N}(SNO)$ are calculated for two values of the threshold (effective) kinetic energy of the final state electron, $T_{\\mathrm{e},\\mathrm{th}} = 6.75$\\ MeV and 5.0 MeV. The published SNO data were obtained using the first value \\cite{SNO1}, while the second one is the threshold energy planned to be reached at a later stage of the experiment.  The {\\it Full Night} D-N asymmetry in the CC event rate in the SNO detector, $A^{N}_{D-N}(SNO)$, can be in the LMA MSW solution region by a factor of $\\sim (1.5 - 2.0)$ bigger than the {\\it Full Night} D-N asymmetry in the solar neutrino induced event rate in the Super-Kamiokande detector \\cite{DNSNO00}: $(A^{N}_{D-N}(SNO))^{LMA} \\cong (1.5 - 2.0) (A^{N}_{D-N}(SK))^{LMA}$. The asymmetry $A^{N}_{D-N}(SNO)$ measured in the SNO experiment can be as large as $(15 - 20)\\%$. A value of $A^{N}_{D-N}(SNO) \\cong 15\\%$, for instance, cannot be excluded by the 95\\% C.L. (2 s.d.) upper limit on $A^{N}_{D-N}(SK)$ following from the Super-Kamiokande data on the D-N effect \\cite{SKsol,SKSmySNO}. In the best fit point of the LMA MSW solution region found in \\cite{FogliSNO,GoswaSNO,ConchaSNO,KSSNO} and in \\cite{ConchaSNO2} we get for $T_{\\mathrm{e},\\mathrm{th}} = 6.75~{\\rm MeV~(~5.0~MeV})$, $(A^{N}_{D-N}(SNO))^{LMA}_{BF1} \\cong 7.3~(6.6)\\%$ and $(A^{N}_{D-N}(SNO))^{LMA}_{BF2} \\cong 10.1~(9.3)\\%$, respectively. At the same time, one finds a considerably smaller value of $A^{N}_{D-N}(SNO)$ in the LMA solution best fit point obtained in \\cite{SKSmySNO}: $(A^{N}_{D-N}(SNO))^{LMA}_{BF3} \\cong 5.0~(4.6)\\%$. In the LMA MSW solution region, the {\\it Core} D-N asymmetry in the SNO detector is predicted to be larger than the {\\it Full Night} D-N asymmetry typically by a factor of $\\sim 1.2$: $(A^{C}_{D-N}(SNO))^{LMA} \\cong 1.2(A^{N}_{D-N}(SNO))^{LMA}$.  In the case of the LOW solution of the solar neutrino problem one has (Figs. 1 and 2) in the region where $A^{N}_{D-N}(SK) > 1\\%$: $A^{N}_{D-N}(SNO) \\cong (1.2 - 1.4)A^{N}_{D-N}(SK)$. In the  solution region given by eq. (5) we find $(A^{N}_{D-N}(SNO))^{LOW} \\cong (1.0 - 7.5)\\%$. In the region under discussion, $(A^{C}_{D-N}(SNO))^{LOW} \\cong (A^{N}_{D-N}(SNO))^{LOW}$. In the best fit point of the LOW solution found in \\cite{FogliSNO,ConchaSNO,GoswaSNO,KSSNO} and in \\cite{ConchaSNO2} we get $(A^{N}_{D-N}(SNO))^{LOW}_{BF1,2} \\cong 3.8~(4.2)\\%$ for $T_{\\mathrm{e},\\mathrm{th}} = 6.75~{\\rm MeV~(~5.0~MeV})$, while in the best fit point obtained in \\cite{SKSmySNO} one has $(A^{N,C}_{D-N}(SNO))^{LOW}_{BF3} \\cong 1.2~(1.5)\\%$. An observation of $A^{N}_{D-N}(SNO) \\gtap 10\\%$ will strongly disfavor the LOW solution of the solar neutrino problem, while an observation of $A^{N}_{D-N}(SNO) > 1\\%$ would rule out the QVO solution.  We have derived also detailed predictions for the ratio of the event rates of the CC reaction $\\nu_e + D \\rightarrow e^{-} + p + p$, $R_{SNO}(CC)$, and of the neutral current (NC) reaction $\\nu + D \\rightarrow \\nu + n + p$, induced by the solar neutrinos in SNO during the {\\it day}, $R^{SNO}_{CC/NC}(D)$, during the {\\it night}, $R^{SNO}_{CC/NC}(N)$, and for the case of the CC event rate produced at night by solar neutrinos which cross the Earth {\\it core}, $R^{SNO}_{CC/NC}(C)$ (Figs. 4 - 6). The predictions were obtained for $T_{\\mathrm{e},\\mathrm{th}} = 6.75~{\\rm MeV~and~5.0~MeV}$. We find that in the LMA MSW solution region given by eqs. (\\ref{dmsolLMA}) and (\\ref{thLMA}), $R^{SNO}_{CC/NC}(X) \\cong (0.20 - 0.65)$, $X=D,N,C$; for $\\Delta m^2 \\ltap 2\\times 10^{-4}~{\\rm eV^2}$ from this region we have $R^{SNO}_{CC/NC}(X) \\cong (0.20 - 0.45)$. In the LOW solution region given by eq. (\\ref{dmsolthLOW}) we obtain $R^{SNO}_{CC/NC}(X) \\cong (0.38 - 0.45)$. In the LMA solution best fit points (see the text) we get $R^{SNO}_{CC/NC}(X) \\cong (0.27 - 0.31)$, while in the two LOW solution best fit points discussed in the text we find approximately $R^{SNO}_{CC/NC}(X) \\cong 0.44~{\\rm and}~0.49$.   In the case of the LMA and LOW solutions, the value of $A^{N}_{D-N}(SNO)$ is very sensitive to the value of $\\Delta m^2$, while $R^{SNO}_{CC/NC}$ exhibits a very strong dependence on $\\tan^2\\theta$. A measured value of $A^{N}_{D-N}(SNO) > 1.0\\%$ and/or of $R^{SNO}_{CC/NC} \\ltap 0.45$ in the SNO experiment can strongly diminish the regions of the allowed values of $\\Delta m^2$ and $\\tan^2\\theta$ of the LMA MSW and of the LOW-QVO solutions of the solar neutrino problem. An upper limit on $A^{N}_{D-N}(SNO)$ in the case of the LMA MSW (LOW-QVO) solution would imply a lower (upper) limit on $\\Delta m^2$. At the same time, an experimental upper limit on $R^{SNO}_{CC/NC}$ would lead to an upper (lower) limit on $\\Delta m^2$. Thus, even upper limits on $A^{N}_{D-N}(SNO)$ of the order of 10\\% and on $R^{SNO}_{CC/NC}$ of the order of 0.45 can significantly reduce the LMA MSW and the LOW-QVO solution regions. \\vspace{-0.5cm}"
    },
    "0201/astro-ph0201447_arXiv.txt": {
        "abstract": "We present a Keck ESI spectrum of FIRST~074711.2+273904, a $K=15.4$ quasar with redshift 4.11 that is detected by both FIRST and 2MASS. The spectrum contains at least 14 independent \\ion{C}{4} absorption systems longward of the Ly$\\alpha$ forest.  These systems are found over a path length of $\\Delta z = 0.984$, constituting one of the highest densities per unit redshift of \\ion{C}{4} absorption ever observed.  One of the \\ion{C}{4} systems is trough-like and resembles a weak BAL-type outflow.  Two of the \\ion{C}{4} are ``associated'' absorption systems with $|v| < 3000\\,{\\rm km\\,s^{-1}}$.  Of the 11 remaining systems with $v > 3000\\,{\\rm km\\,s^{-1}}$, eight are either resolved or require multiple discrete systems to fit the line profiles.  In addition to \\ion{C}{4} absorption, there are two low-ionization \\ion{Mg}{2} absorption systems along with two damped Ly$\\alpha$ systems, at least one of which may be a \\ion{C}{4} system. The overdensity of \\ion{C}{4} absorption spans a redshift range of $\\Delta z \\sim 1$.  Superclusters along the line of sight are unlikely to cause an overdensity stretching over such a long redshift path, thus the absorption may be an example of narrow, high-velocity, intrinsic absorption that originates from the quasar.  We suggest that this quasar is a member of a transitional class of BAL quasars where we are just barely seeing the spatial, density, or temporal edge of the BAL-producing region (or period); the multiple high-velocity absorption systems may be the remnants (or precursors) of a stronger BAL outflow.  If correct, then some simpler absorption line complexes in other quasars may also be due to outflowing rather than intervening material. ",
        "introduction": "It is generally agreed that the broad absorption line (BAL) systems seen in the spectra of many quasars are intrinsic to the quasar, arising from high velocity outflow of gas directly from the accretion disk region \\citep{wmf+91}.  Narrow absorption lines in quasar spectra, however, arise from a variety of sources.  Narrow absorption lines with small velocities relative to the quasar emission redshift could be caused by galaxies along the line of sight, clouds in the interstellar medium of the host galaxy, or even smaller scale gas flows within a few parsecs of the black hole.  On the other hand, most of the narrow absorption systems that have large velocities with respect to the emission redshift of the quasar are thought to be intervening \\citep{ssb88}; however, there is increasing evidence that even some fraction of these may be intrinsic.  In many BAL quasars, the absorption troughs have considerable structure, suggesting a connection to narrow absorption line systems.  We discuss the spectrum of a newly discovered quasar, FIRST J074711.2+273904 (hereafter FIRST~0747+2739), with an overdensity of \\ion{C}{4} absorbers in its spectrum.  It may be an example of an object whose BAL component has been caught in a transitional state that could be related to orientation, time, or density. ",
        "conclusions": "In summary, we find that that overdensity of \\ion{C}{4} absorption lines seen along the line of sight to FIRST~0747+2739 is sufficiently high that it is likely that at least some of these systems are high-velocity, intrinsic absorption systems.  The presence of weak trough-like absorption features suggests that FIRST~0747+2739 may be a transitional BAL quasar that is being observed at a brief phase in its transition to (or from) a standard quasar, or is being viewed at a highly specific orientation, or that it is a BAL with a very low absorber density.  If the absorption is intrinsic, monitoring with high resolution spectroscopy over the course of $\\gtrsim 18$ months may reveal variability of the absorbers, a signature of intrinsic absorption.  Confirmation (or exclusion) of the intrinsic hypothesis for the unusually large number of high-velocity \\ion{C}{4} systems in FIRST~0747+2739 would have interesting consequences for \\ion{C}{4} absorption line studies, particularly high-redshift superclustering studies."
    },
    "0201/astro-ph0201392_arXiv.txt": {
        "abstract": "The Wisconsin H-Alpha Mapper (WHAM) is a high throughput Fabry-Perot facility developed specifically to detect and explore the warm, ionized component of the interstellar medium at high spectral resolution.  It began operating at Kitt Peak, Arizona in 1997 and has recently completed the WHAM Northern Sky Survey (WHAM-NSS), providing the first global view of the distribution and kinematics of the warm, diffuse H~II in the Milky Way.  This H$\\alpha$ survey reveals a complex spatial and kinematic structure in the warm ionized medium and provides a foundation for studies of the temperature and ionization state of the gas, the spectrum and strength of the ionizing radiation, and its relationship to other components of the interstellar medium and sources of ionization and heating within the Galactic disk and halo. ",
        "introduction": "Warm ionized gas is a principal component of the interstellar medium in our Galaxy and others.  Its large scale height, mass surface density, and power requirement have significantly modified our understanding of the composition and structure of the interstellar medium and the distribution and flux of ionizing radiation within the disk and halo (e.g., Kulkarni \\& Heiles 1987; McKee 1990, Reynolds 1991, Ferri\\`{e}re 2001).  Although originally detected in the 1960s with radio techniques, subsequent developments in high-throughput Fabry-Perot spectroscopy have shown that the primary source of information about the distribution, kinematics, and other physical properties of this gas is obtained through the detection and study of faint, diffuse interstellar emission lines at optical wavelengths.  Presented below are some recent results from the Wisconsin H$\\alpha$ Mapper (WHAM), including velocity-interval maps from the recently completed WHAM Northern Sky Survey (WHAM-NSS) of interstellar H$\\alpha$ as well as observations of much fainter ``diagnostic'' emission lines that probe the ionization and excitation state of the gas.  \\subsection{The Warm Ionized Medium in the Milky Way}  Diffuse ionized gas is a major, yet poorly understood component of the interstellar medium, which consists of regions of warm (10$^{4}$ K), low-density (10$^{-1}$ cm$^{-3}$), nearly fully ionized hydrogen that occupy approximately 20\\% of the volume within a 2 kpc thick layer about the Galactic midplane (e.g., Haffner, Reynolds, \\& Tufte 1999).  Near the midplane, the space averaged density of H~II is less than 5\\% that of the H~I. However, because of its greater scale height, the total column density of interstellar H~II along high Galactic latitude sight lines is relatively large, 1/4 to 1/2 that of the H~I, and one kiloparsec above the midplane, warm H~II may be the dominant state of the interstellar medium (Ferri\\`ere 2001; Reynolds 1991b).  The presence of this ionized medium can have a significant effect upon the interstellar pressure near the Galactic midplane (Cox 1989) and upon the dynamics of hot (10$^5$ -- 10$^6$ K), ``coronal'' gas far above the midplane (e.g., Heiles 1990). Miller \\& Cox (1993) have suggested that this gas is part of a wide spread, warm intercloud medium, while in the McKee \\& Ostriker (1977) picture of the interstellar medium, this warm H~II is located in the outer envelopes of H~I clouds, forming the boundary between the clouds and a wide spread, hot (``coronal'') phase.  It is generally believed that the O stars, confined primarily to widely separated stellar associations near the Galactic midplane, are somehow able to account for this widespread gas, not only in the disk but also within the halo, 1-2 kpc above the midplane.  However, the nature of such a disk-halo connection is not clear.  For example, the need to have a large fraction of the Lyman continuum photons from O stars travel hundreds of parsecs through the disk seems to conflict with the traditional picture of H~I permeating much of the interstellar volume near the Galactic plane. It has been suggested that ``superbubbles'' of hot gas, especially superbubbles that blow out of the disk (``galactic chimneys''), may sweep large regions of the disk clear of H~I, allowing ionizing photons from the O stars within them to travel unimpeded across these cavities and into the halo (e.g., Norman 1991).  Another possibility is that the Lyman continuum radiation itself is able to carve out extensive regions of H~II through low density portions of the H~I (e.g., Miller and Cox 1993), perhaps creating photoionized pathways or ``warm H~II chimneys'' that extend far above the midplane (Dove and Shull 1994; Dove, Shull, and Ferrara 2000). Although the existence of superbubbles has long been established (e.g., Heiles 1984), direct observational evidence that such cavities are actually responsible for the transport of hot gas and ionizing radiation up into the Galactic halo is very limited.  Interestingly, even though the source of ionization is believed to be O stars, the temperature and ionization conditions within the diffuse ionized gas appear to differ significantly from conditions within classical O star H~II regions.  For example, anomalously strong [S~II] $\\lambda$6716/H$\\alpha$ and [N~II] $\\lambda$6584/H$\\alpha$, and weak [O~III] $\\lambda$5007/H$\\alpha$ emission line ratios (compared to the bright, classical H~II regions) indicate a low state of excitation with few ions present that require ionization energies greater than 23 eV (Haffner et al 1999; Rand 1997).  This is consistent with the low ionization fraction of helium, at least for the helium near the midplane (Reynolds \\& Tufte 1995; Tufte 1997; Heiles et al 1996), which implies that the spectrum of the diffuse interstellar radiation field that ionizes the hydrogen is significantly softer than that from the average Galactic O star population.  Rand (1997) has also reported lower helium ionization in the H~II halo of the edge-on galaxy NGC 891.  Furthermore, it has recently become apparent that O star photoionization models fail to explain observed spatial variations in some of the line intensity ratios.  For example, the models do not explain the very large increases in [N~II]/H$\\alpha$ and [S~II]/H$\\alpha$ (accompanied by an increase in [O~III]) with distance from the midplane or the observed constancy of [S~II]/[N~II] (see discussions by Reynolds et al 1999, Haffner et al 1999, and Collins \\& Rand 2001). The data seem to require the existence of a significant \\emph{non-ionizing} source of heat that overwhelms photoionization heating at low densities within the ionized medium (Reynolds et al 1999; Collins \\& Rand 2001; Otte et al 2001; Bland-Hawthorn, Freeman, \\& Quinn 1997).  Proposed sources include the dissipation of MHD turbulence, Coulomb interactions with cosmic rays, magnetic reconnection, and photoelectric heating by a population of very small grains (see Minter \\& Spangler 1996; Reynolds et al 1999; Weingartner \\& Draine 2001).   \\subsection{WHAM}  The Wisconsin H-Alpha Mapper (WHAM) is a remotely controlled observing facility, funded by the National Science Foundation and dedicated to the detection and study of faint optical emission lines from the diffuse ionized gas in the disk and halo of the Milky Way (Tufte 1997; Reynolds et al 1998b, Haffner 1998).  The WHAM facility consists of a 15 cm aperture dual-etalon Fabry-Perot spectrometer (the largest used in astronomy) coupled to a 0.6 m aperture siderostat, which provide a one-degree diameter beam on the sky and produce a 12 km s$^{-1}$ resolution spectrum across a 200 km s$^{-1}$ spectral window.  The spectral window can be centered on any wavelength between 4800 \\AA\\ and 7300 \\AA\\ using a gas (SF$_6$) pressure (optical index) control system and a filter wheel.  The tandem etalons greatly extend the effective ``free spectral range'' of the spectrometer, improve the shape of the response profile, and suppress the multi-order Fabry-Perot ghosts, especially those arising from the relatively bright atmospheric OH emission lines within the pass band of the interference filter.  A high quantum efficiency (78\\% at H$\\alpha$), low noise (3 e$^{-}$ rms) CCD camera serves as a multichannel detector, recording the spectrum as a Fabry-Perot ``ring image'' without scanning (e.g., Reynolds et al 1998b).  The construction and testing of this facility at the University of Wisconsin was completed in September 1996, and WHAM began operating on Kitt Peak in January 1997 (see Fig. 2). Since then, WHAM has been successfully collecting data nearly every clear, dark-of-the-moon period. It has completed as its first major mission a 37,565 spectra H$\\alpha$ survey of the northern sky, which has provided the first map of the large scale distribution and kinematics of diffuse interstellar H~II that is comparable to earlier 21~cm surveys of H~I (\\S 2, below). WHAM is now beginning its second major mission, a comprehensive study of fainter, diagnostic emission lines that trace the excitation and ionization conditions within the gas as well as the strength and spectrum of the ionizing radiation (\\S 3 \\& \\S4, below). ",
        "conclusions": "The development of high throughput, high spectral resolution Fabry-Perot spectroscopy has established that faint, diffuse interstellar emission lines at optical wavelengths contain a wealth of new information about the interstellar medium that cannot be obtained through other techniques at other wavelengths.  The WHAM H$\\alpha$ survey plus the detection of fainter nebular lines reveal a complexity not only in the structure and kinematics of the warm ionized medium, but also in its excitation and ionization conditions.  Studies of this weak emission have begun to shed new light on the nature of the interstellar medium and the principal sources of ionization and heating within the disk and halo of the Galaxy."
    },
    "0201/astro-ph0201422_arXiv.txt": {
        "abstract": "We report our first results on comparing the variations of the solar internal rotation with solar activity, as predicted by non-linear solar dynamo modelling, with helioseismic measurements using the SOHO MDI data. ",
        "introduction": "The migrating bands of faster and slower rotation of the solar surface were first discovered by Howard and LaBonte (1980). By analyzing helioseismic data which are now provided by the space (SOHO MDI) and ground-based (GONG) projects, Howe et al (2000) and Antia and Basu (2000) have found that these ``torsional oscillations'' penetrate quite deep into the solar interior, to at least 8 percent of solar radius.  The mechanism responsible for producing the 11-yr solar torsional oscillations is thought to be the non-linear interaction between the magnetic field and the solar differential rotation. Comparing the spatial and temporal structure of the torsional oscillations predicted by the theoretical modelling with helioseismic measurements would allow the calibration of the theoretical models of the solar dynamo, leading finally to better understanding of the basic mechanisms of solar magnetic activity.  In this contribution, we address the predictions of a two-dimensional axisymmetric mean-field dynamo model in a spherical shell, in which the only nonlinearity is the action of the azimuthal component of the Lorentz force of the dynamo-generated magnetic field on the solar angular velocity (Covas et al 2000). The torsional oscillations produced in this model are compared with the results of helioseismic inversion of the SOHO MDI data, now available over almost half of the 11-yr solar activity cycle.  Helioseismic measurements are based on analyzing the rotational splittings of the solar p-mode frequencies. These have different sensitivities to the rotation at different depths and latitudes, which can distort the actual rotation profiles when they are inverted from the data of finite accuracy. We address this problem by using artificial inversions. ",
        "conclusions": "When comparing the torsional oscillations predicted by the theoretical modelling with helioseismic data, we observe some general features which are in common, as well as significant differences.  Both the model and the observations show that all the convection zone, down to its base, is involved in the oscillations. Preliminary experiments suggest that this feature remains even when density stratification is included. The ``zonal flows'' propagate towards the equator at lower latitudes and towards the pole at higher latitudes, and the model flows show the correct phasing with solar activity.  The observed low-latitude zonal flow, however, is much more localized in depth, compared with model prediction, and closer to the surface. The high-latitude accelerating flow is much stronger than in the model, has a larger latitudinal extent, and is situated at somewhat lower latitudes.  The differences are hardly surprising; indeed, in this very first comparison we made no attempts to tune the model parameters to fit the observations, as our interest was in the effect of noise on the inversions of a known data set. These differences are  also related to the inherent uncertainties in, and simplified nature of, the dynamo model: for example uniform density is assumed.  With the current accuracy of the seismic data, we are not able to resolve the well-structured variations near the base of the convection zone, predicted by the dynamo model. We deduce that such structures, if present in the `real' Sun, also would not be reliably detected by current techniques. We believe that the accuracy of the seismic measurements will improve significantly in the near future -- due to improved accuracy in the measurement of the rotational splittings in the solar data when using the  more sophisticated techniques which are being developed, and also from use of the data from ground-based observations, together with better coverage of the solar cycle.  Even with the accuracy which is now available, the seismic measurements of the torsional oscillations provide new and valuable constraints on the physical modelling of the solar dynamo. We believe that we have also shown that dynamo models can have an input into the interpretation of the seismic data."
    },
    "0201/astro-ph0201087_arXiv.txt": {
        "abstract": "Presented here are high angular resolution MERLIN 5 GHz (6 cm) continuum observations of  the binary proplyd system, LV 1 in the Orion nebula, which consists of proplyd 168--326SE and its binary proplyd companion  168--326NW (separation $0.4\\arcsec$). Accurate astrometric alignment allows a detailed comparison between these data and published \\emph{HST} PC H$\\alpha$ and \\oiii\\ images.  Thermal radio sources coincide with the two proplyds and originate in the ionized photoevaporating flows seen in the optical emission lines. Flow velocities of $\\approx\\ 50 \\kms$ from the ionized proplyd surfaces and $\\geq 100 \\kms$ from a possible micro-jet have been detected using  the Manchester Echelle spectrometer.  A third radio source is found to coincide with a region of extended, high excitation,  optical line emission that lies between the binary proplyds  168--326SE/326NW . This is modelled as a bowshock due to the collision of the photoevaporating flows from the two proplyds. Both a thermal and a non-thermal origin for the radio emission in this collision zone are considered. ",
        "introduction": "It is now well established that the compact, highly ionized gaseous knots (LV 1-6) near the Trapezium stars in M42, discovered by Laques \\& Vidal (1979) contain young low mass stars (YSOs) still partially cocooned in their primaeval material. See Figure 1 for their locations.  Many similar objects were found in the same vicinity at radio wavelengths (Churchwell et al. 1987; Garay 1987; Garay, Moran, \\& Reid 1987; Felli et al. 1993a; Felli et al. 1993b). The LV knots are particularly interesting  both because of their close proximity to the Sun, permitting detailed observation, and because of their extreme local environment. This environment is swept by the energetic particle wind of the nearby O6 star, \\thetc\\ and the LV knots are irradiated by its Lyman (and near UV) photons. Their outside surfaces are consequently highly ionized, though shielded form the direct wind from \\thetc\\ by a standoff bowshock (Figure 3) in the photoevaporated flow.  \\emph{HST} imagery revealed the structure of these systems, now known as `proplyds' (O'Dell, Wen, \\& Hu 1993), with startling clarity (O'Dell \\& Wong 1996; Bally et al. 1998). The central stars, found originally on a 2 $\\mu$m IRCAM `engineering' image with UKIRT by Meaburn (1988) and subsequently by McCaughrean \\& Stauffer (1994) are seen to be embedded in dense  cocoons/disks of dusty  molecular gas. Photoevaporated flows from the ionized surfaces form cometary tails pointing away from \\thetc\\ as they meet the UV flux of this star. Supersonic flow velocities of $\\geq$ 50\\kms\\ were found (Meaburn 1988; Meaburn et al. 1993; Massey \\& Meaburn 1995) using the Manchester Echelle spectrometer (MES - Meaburn et al. 1984) on a variety of telescopes. The origin of these ionized flows has been modelled, initially as a two-wind interaction (Henney et al. 1996, 1997) and more recently as the interaction of the flow from the proplyd disk with the ionizing Lyman photon field of \\thetc\\ and the nebula (Johnstone, Hollenbach, \\& Bally 1998; Henney \\& Arthur 1998; Henney \\& O'Dell 1999; St{\\\"o}rzer \\& Hollenbach 1999). The two-wind interaction is now thought to manifest itself much further from the proplyd disk in the form of the standoff bowshocks (one of which is shown in Figures 1 and 3) observed by Hayward, Houck and Miles (1994),   McCullough et al. (1995) and  Bally et al. (1998). The flow from the proplyd itself is most likely caused  by photoevaporation of the disk by either ionizing (EUV) or non-ionizing (FUV) photons from \\thetc, depending on the proplyd's proximity to it (Johnstone, Hollenbach, \\& Bally 1998). The process of photoevaporation, photoionization, and acceleration of the disk gas down the surrounding density gradient is found to produce flow velocities of only 2-3 times the sound speed of $\\sim$10\\kms\\ (Dyson 1968), which does not fully account for the velocities observed in the longslit spectra. These longslit spectra were complemented by the Fabry-Perot work of O'Dell et al. (1997), which compared velocity data with \\emph{HST} images and suggested the presence of a micro-jet associated with one of the proplyds.  The presence of predominantly monopolar jets from the YSOs in many of these proplyds was first suggested by the $\\geq$ 100\\kms, spatially compact ($\\leq$ 1 arcsec across) spikes or knots on the MES position-velocity (pv) arrays of Meaburn et al. (1993) and Massey \\& Meaburn (1995).  This has now been confirmed through direct \\emph{HST} imagery by Bally, O'Dell, \\& McCaughrean (2000), who find more than 20 jets associated with various proplyds.  This apparent ubiquity of jets and the fact that it is difficult to get high velocity flows from density gradients suggest that most collimated high velocity features are jets.  \\emph{HST} imagery, as a consequence of its high angular resolution, along with high resolution near infra-red imagery (Petr 1998),  has also revealed the binary nature of a small fraction of the proplyds.  The binary structure of the prominent proplyd originally designated LV 1 when unresolved in ground-based observations is of particular interest. It was first resolved into two sources in the 2 cm VLA A-array radio maps of Felli et al. (1993a) (Sources 5 and 20 in their Tables 2 and 3).  The higher resolution \\emph{HST} image shows that this system is formed by two proplyds separated by 0.4\\arcsec, which are catalogued as 168--326N and as 168--326S, by O'Dell \\& Wen (1994). In the subsequent full catalogue of stars and compact objects of O'Dell \\& Wong (1996), it appears as a single proplyd, 168--326, corresponding to Felli et al.'s (1993a) source 20.  However, the near infra-red 0.13\\arcsec\\ resolution binary star survey of Petr (1998) identifies 168--326 as a binary system consisting of two components that are referred to as 168--326E and 168--326W.  Bally et al. (1998), who used the same \\emph{HST} observations as used here, label the two components as 168--326 and 168--326a (their figure 2) and mention the presence of a bright arc lying between the sources.  Here, the proplyds are referred to as 168--326SE and 168--326NW, which is the clearest form of designation.  In the present paper the discovery is reported of the interaction zone between the flows from 168--326SE and 326NW in both the radio (at 5 GHz with MERLIN) and optical (\\ha\\ and \\oiii\\ with the \\emph{HST}) domains. The collision of two photoevaporated flows is explored to explain the nature of this interaction.  Further kinematical and morphological evidence for a jet in LV 1 is also presented. ",
        "conclusions": "The Orion proplyd originally designated LV 1 is shown at both radio (with MERLIN) and optical wavelengths (with the HST) to be a proplyd binary system with an interaction zone separating the two components 168-326NW and 168-326SE.  Comparison of detailed hydrodynamical modelling of the interaction zone with the observations has identified the geometry of the system convincingly.  A strong 6 cm radio source coincides with the stagnation point in this interaction zone. The present observations cannot distinguish between a thermal or non-thermal origin for this radio emission. Observations at 20 cm but with a similar angular resolution (0.1 \\arcsec) are awaited to clarify this point through measurement of the radio spectral index.  The existence of a  high-speed micro-jet from the binary component 168-326SE is also suggested by spectral observations and HST imagery.  \\vspace{10mm}  \\noindent {\\bf Acknowledgements}  MFG would like to acknowledge a studentship received from the Particle Physics and Astronomy Research Council (PPARC).  MERLIN is a national facility operated by the University of Manchester on behalf of PPARC in the U.K. The work presented in this paper is based on observations made with MERLIN and the NASA/ESA Hubble Space Telescope, the latter obtained from the data archive at the Space Telescope Science Institute. STScI is operated by the Association of Universities for Research in Astronomy, Inc. under NASA contract NAS 5-26555."
    },
    "0201/astro-ph0201278_arXiv.txt": {
        "abstract": "We present new 37.7\\,$\\mu$m far-infrared imaging of the infrared luminous (L$_{\\rm IR}\\sim 5.16\\times10^{11}$L$_\\sun$) interacting galaxy Arp\\,299 (= IC\\,694 + NGC\\,3690). We show that the 38\\,$\\mu$m flux, like the 60 and 100\\,$\\mu$m emission, traces the luminosity of star forming galaxies, but at considerably higher spatial resolution. Our data establish that the major star formation activity of the galaxy originates from a point source in its eastern component, IC\\,694, which is inconspicuous in the optical, becoming visible only at the near and mid-infrared. We find that IC\\,694 is two times more luminous than NGC\\,3690, contributing to more than 46\\% of the total energy output of the system at this wavelength. The spectral energy distribution of the different components of the system clearly shows that IC\\,694, has 6 times the infrared luminosity of M82 and it is the primary source responsible for the bolometric luminosity of Arp\\,299. ",
        "introduction": "One of the major results of the IRAS survey was the discovery of a class of extremely luminous galaxies (L$\\ge$10$^{11}\\,$L$_\\sun$) which emit most of their energy in the far-infrared \\citep{Houck84,Soifer89}.  These ultraluminous infrared galaxies (ULIRGs) have been intensely studied over the past 15 years leading to significant progress in our understanding of their properties \\citep[see review of][]{Sanders96}. For example, it is now clearly established that most, if not all, ULIRGs are interacting systems \\citep[i.e.][]{Clements96,Duc97,Borne00}, with ample quantities of molecular gas and dust surrounding their active central regions \\citep{Solomon97,Gao99}. It appears likely that the mergers caused the inordinate far-infrared luminosities, either by compressing the natal ISM and triggering a global starburst, or by triggering accretion onto a central super massive black hole forming an AGN \\citep{Genzel98,Sanders99,Joseph99}. It is not clear, however, which mechanism dominates.  This is because massive star formation is always expected near the central potential well of ULIRGs, but the spatial resolution at mid- to far-infrared wavelengths, where most of the energy of those galaxies appears, is rather poor.  For example, for Arp\\,220, the prototypical ULIRG, its two putative nuclei are separated by $\\sim1''$, and most of the molecular gas is well within the central 10$''$ \\citep{Scoville91}. The peak of the far-infrared spectral energy distribution for Arp 220 though lies near 60\\,$\\mu$m where the spatial resolution available from IRAS and ISO is greater than 30$''$ -- insufficient to resolve the locations of the heating source(s). Consequently we rely on less direct diagnostics, usually in the near and mid-infrared, to address the questions of the origins of the far-infrared luminosity \\citep{Veilleux95,Genzel98,Murphy99,Laurent00,Soifer00}. This approach has proven to be extremely powerful and has produced many surprises. One prime example was the mid-infrared imaging of the Antennae galaxies (NGC\\,4038/39) with the Infrared Space Observatory which revealed that 15\\% of its energy at 15\\,$\\mu$m originated from a star cluster inconspicuous in the optical and near-infrared \\citep{Mirabel98}. However, while observations in the mid-infrared can much better probe the dusty cores of IR bright galaxies than visible observations, the cores can still be extincted in the mid-infrared.  Furthermore, only a small fraction ($\\sim3$\\%) of the total luminosity of luminous infrared galaxies emerges in the mid-infrared \\citep{Laurent00}, so that we are still not directly probing the regions where the far-infrared flux originates\\footnote{For normal late type galaxies it has been shown that $\\sim$15\\% of the luminosity is emitted between 5--20$\\mu$m \\citep{Dale01}}.  Arp\\,299\\,(=Mrk\\,171, VV\\,118, IC\\,694+NGC\\,3690), at a distance of 41\\,Mpc, (v$_{\\rm hel}$ = 3080\\,km\\,s$^{-1}$, H$_{\\rm 0}$=75\\,km\\,s$^{-1}$\\,Mpc$^{-1}$), is a relatively nearby peculiar galaxy. The system is dynamically young, its components still violently interacting leading to the formation of an 180\\,kpc (13$'$) long tidal tail \\citep{Hibbard99}. Diffuse emission around the central regions extends over an area 1.5$'$ in size, while the galaxy itself consists of IC\\,694\\,(=Mrk171B = UGC6472 = VV118a) to the east and NGC\\,3690\\,(=Mrk171A = UGC6471 = VV118b) nearly 20$''$ to the west (see Figure~\\ref{image}a). Following the nomenclature of \\citet{Gehrz83} and \\citet{ww91}, the nucleus of IC\\,694 is often called source A.  NGC\\,3690 is resolved to sources B1 and B2 to the south, with B2 marking the position of its nucleus, and sources C and C$'$ located $\\sim7''$ to the north of B2. It is believed that Arp\\,299 results from a prograde-retrograde encounter between a gas rich Sab-Sb galaxy (IC\\,694) and an SBc-Sc galaxy (NGC\\,3690) that occured 750\\,Myrs ago, and that the system will merge in 20--60\\,Myrs \\citep{Hibbard99}. Sources C and C$'$ do not seem to have significant potential wells to be considered individual galaxies given the lack of evidence for an underlying concentration of old, red stars \\cite[i.e.][]{Alonso00}. Due to its proximity, Arp\\,299 has been studied extensively ever since it was discovered to have the most luminous emission lines of any non-Seyfert Markarian galaxy \\citep{Weedman72}. It is among the most X-ray luminous galaxies \\citep[L$_{\\rm X}\\sim10^{42}$\\,ergs\\,s$^{-1}$][]{Fabbiano92,Zezas98} and it is also infrared luminous since, based on its IRAS faint source catalogue fluxes \\citep{Moshir90}, one can calculate\\footnote{We use the standard definition of L$_{\\rm IR}$(8--1000\\,$\\mu$m)=5.62$\\times$10$^5$\\,D(Mpc)$^2$\\,(13.48{\\it f}$_{12}$+5.16{\\it f}$_{25}$+2.58{\\it f}$_{60}$+{\\it f}$_{100}$)\\,L$_\\sun$ \\citep[see][]{Sanders96}.}  that L$_{\\rm IR}=5.16\\times10^{11}$L$_\\sun$. Large quantities of molecular gas with strong streaming motions are observed in the system, providing ample fuel for the star formation activity taking place \\citep{Sargent91,Aalto97,Casoli99}. Recent high resolution imaging has revealed that the morphology of NGC\\,3690, which is brighter than IC\\,694 in the optical and near infrared, is quite complex \\citep{Lai99,Alonso00}. Even though there is no compelling evidence for an active galactic nucleus (AGN) anywhere in the system \\citep[i.e.][]{Augarde85,Armus89,Smith96} it is clear that the extinction due to dust is high, considerably changing the apparent morphology of the galaxy as one observes it from the UV and optical to near- and mid-infrared wavelengths \\citep{Gehrz83,Joy89,Dudley93,Gallais99,Dudley99,Xu00,Soifer01}.  In the present paper we will focus our attention on IC\\,694. While this component appears quite diffuse in the UV and optical, its bulge becomes obvious as a point source only at near-infrared \\citep{ww91,Lai99}. The description of our observations is presented in Section 2, the implications of our findings to the current understanding of Arp\\,299 are discussed in Section 3, and finally our conclusions are summarized in Section 4. ",
        "conclusions": "We have obtained high resolution 38\\,$\\mu$m mid-infrared images of the peculiar galaxy Arp\\,299. Our data clearly shows that despite its diffuse and quiescent appearance in the UV and optical, IC\\,694 harbors a strong point-like source in the mid and far-infrared and is the dominant source of far-infrared radiation in the system.  Together with shorter wavelength images, we construct the spectral energy distribution of the various components of the galaxy.  We find that IC\\,694 is by far the strongest source, with an infrared luminosity of 1.8$\\times10^{11}$\\,L$_\\sun$, or $\\sim$40\\% of the whole Arp\\,299 system.  The infrared luminosity of IC\\,694 is 6 times the luminosity of M\\,82, while the one inferred for component C is about 1.5 times that of M\\,82, making it {\\em one of the most luminous non-nuclear starbursts known}.  Our analysis suggests that in order to accurately determine the starburst activity and infrared luminosity of different regions in interacting/merging luminous infrared galaxies it is imperative to obtain good spatial resolution maps covering the 15--40\\,$\\mu$m range to better trace the colder dust component. Future instruments such as IRS and MIPS on SIRTF and FORCAST on SOFIA will provide valuable information in addressing these problems."
    },
    "0201/astro-ph0201034_arXiv.txt": {
        "abstract": "The evidence for positive cosmological constant $\\Lambda$ from Type Ia supernovae is reexamined.  Both high redshift supernova teams are found to underestimate the effects of host galaxy extinction. The evidence for an absolute magnitude- decay time relation is much weakened if supernovae not observed before maximum light are excluded. Inclusion of such objects artificially supresses the scatter about the mean relation.  With a consistent treatment of host galaxy extinction and elimination of supernovae not observed before maximum, the evidence for a positive lambda is not very significant (3-4 $\\sigma$). A factor which may contribute to apparent faintness of high z supernovae is evolution of the host galaxy extinction with z.  The Hubble diagram using all high z distance estimates, including SZ clusters and gravitational lens time-delay estimates, does not appear inconsistent with an $\\Omega_o$ = 1 model.  Although a positive $\\Lambda$ can provide an, albeit physically unmotivated, resolution of the low curvature implied by CMB experiments and evidence that $\\Omega_o <$ 1 from large-scale structure, the direct evidence from Type Ia supernovae seems at present to be inconclusive. ",
        "introduction": "The claim that the measured brightnesses of Type Ia supernovae at redshifts 0.1 - 1.0 imply $\\Lambda > 0$ (Schmidt et al 1998, Garnavich et al 1998, Riess et al 1998, Perlmutter et al 1999, Fillipenko and Riess 2000, Riess et al 2001, Turner and Riess 2001) has had a dramatic effect on cosmology.  The model with $\\lambda_o$ = 0.7, where $\\lambda_o = \\Lambda/3 H_o^2$, and $\\Omega_0$ = 0.3 has become a concensus model, believed to be consistent with most evidence from large-scale stucture and CMB fluctuations.  In this paper I test the strength of the evidence that $\\Lambda > 0$ and show that there are inconsistencies in the way the supernovae data have been analyzed.  When these are removed, the strength of the evidence for $\\Lambda > 0$ is much diminished.  To set the scene, Fig 1 shows $B_{max}$ versus log (cz)  for 117 Type Ia supernovae since 1956 from the Barbon et al (1998) catalogue (excluding those labelled '*' which are discovery magnitudes only), together with published supernovae from the high z programmes, corrected for Galactic and internal extinction, but not for decay-time effects, together with predicted curves from an $\\Omega_o$ = 1 model.  At first sight there is not an enormous difference between the high z and low z supernovae, except that the latter seem to show a larger scatter.  Fig 2 shows the same excluding less reliable data (flagged ':', in the Barbon et al catalogue, or objects with pg magnitudes only (Leibundgut et al (1991)), correcting for peculiar velocity effects (see section 3), using the Phillips et al (1999) internal extinction correction (see section 2) where available, and deleting two objects for which the dust correction is $>$ 1.4 mag.  The scatter for the low z supernovae appears to have been reduced.  Finally Fig 3 shows the supernovae actually used by Perlmutter et al (1999).  Now the scatter in the low z supernovae is not much different from the high z supernovae and a difference in absolute magnitude between low z and high z supernovae, relative to an $\\Omega_o$ = 1 model, can be perceived.  However comparison with Fig 2 suggests that the low z supernovae used may be an abnormally luminous subset of all supernovae. We will return to this point in section 4.  Excellent recent reviews of Type Ia supernovae, which fully discuss whether they can be thought of as a homogenous population, have been given by Branch (1998), Hildebrand and Niemeyer (2000) and Leibundgut (2000, 2001).  In this paper I shall assume that they form a single population and that their absolute magnitude at maximum light depends, at most, on a small number of parameters.  I do not, for example, consider the possibility of evolution, discussed by Drell et al (2000).  Absolute magnitudes are quoted for $H_o$ = 100 throughout.  \\begin{figure} \\epsfig{file=snplotnopec2.ps,angle=0,width=8cm} \\caption{ $B_{max}$ versus log (cz) for all Type Ia supernovae post 1956. Magnitudes are corrected for Galactic extinction, and for internal extinction (using the de Vaucouleurs prescription - see section 2).  A mean (B-V) colour of 0.0 has been assumed for the supernovae for which only V at maximum light is known.  Solid curves are loci for $\\Omega_o = 1$ model, with $M_B = -19.64, -17.44 (H_o=100)$, corresponding to $<M_B> \\pm 2 \\sigma$ for solution (1) of Table 2.} \\end{figure}  \\begin{figure} \\epsfig{file=snplotcorr2.ps,angle=0,width=8cm} \\caption{ $B_{max}$ versus log (cz) for Type Ia supernovae post 1956, excluding less reliable data, correcting for effect of peculiar velocity, and using the Phillips et al (1999) internal extinction correction where available. Solid curves as for Fig 1.} \\end{figure}  \\begin{figure} \\epsfig{file=snplot2.ps,angle=0,width=8cm} \\caption{ $B_{max}$ versus log (cz) for Type Ia supernovae used by Perlmutter et al (1999).  No correction for internal extinction is applied. Solid curves are as for Fig 1, but shifted by +0.33 mag. to account for non-correction for extinction.} \\end{figure} ",
        "conclusions": "(1) I have reanalyzed the evidence that high-z supernovae support a universe with positive $\\Lambda$.  (2) Both high-z supernova teams appear to have underestimated host galaxy extinction.  (3) The evidence for an $M_B - \\Delta m_{15}$ relation is weaker than previously stated (only 2.6 $\\sigma$) if analysis is restricted to supernovae observed before maximum.  The rms deviation about the mean relation is significantly larger than previously claimed.  (4) After consistent corrections for extinction are applied the significance of the difference in absolute magnitude between high and low z supernovae, in an Einstein de Sitter ($\\Omega_o$ = 1) universe, is 2.8-4.6 $\\sigma$, depending whether (and how) the $M_B - \\Delta M_{15}$ correction is applied, so such a model can not really be rejected conclusively by the present data.  (5) The Hubble diagram based on all high redshift estimates supports an Einstein de Sitter universe. The HDF-N supernova favours such a universe also, contrary to the published claims of Riess et al (2001).  (6) The community may have been too hasty in its acceptance of a positive $\\Lambda$ universe, for which no physical motivation exists, and needs to reconsider the astrophysical implications of the more natural Einstein de Sitter, $\\Omega_o$ =1, model.  For the supernova method, the need is to continue study of low z supernovae to improve understanding of extinction and of the absolute magnitude decay-time relation, and to consider shifting towards infrared wavelengths, as advocated by Meikle (2000), in order to reduce the effects of extinction.  Of course the arguments presented here do not prove that $\\Lambda$ = 0.  The combination of the evidence from CMB fluctuations for a spatially flat universe with a variety of large-scale structure arguments for $\\Omega_o$ = 0.3-0.5 may still make positive $\\Lambda$ models worth pursuing.  However it would seem to be premature to abandon consideration of other alternatives."
    },
    "0201/astro-ph0201202_arXiv.txt": {
        "abstract": "Why do the circumnuclear (inner 1--2\\,kpc) regions of spirals show vastly different star formation rates (SFR) even if they have a comparable molecular gas content? Why do some develop starbursts  which are intense short-lived (t\\,$\\ll$\\,1\\,Gyr) episodes of star formation  characterized by a high star formation rate per  unit mass of molecular gas (SFR/M$_{\\mathrm{H2}}$), which I refer to as  star formation efficiency (SFE)\\@. I address these questions  using high resolution (2$''$ or 100--200\\,pc) CO  (J=1$\\rightarrow$0) observations from the Owens Valley Radio Observatory, optical  and NIR  images, along with published radio continuum (RC)  and Br$\\gamma$ data. The sample of eleven galaxies  includes  the brightest nearby starbursts comparable to M82 and control non-starbursts. More detailed results are in \\cite{J99} and \\cite{JKS01}. ",
        "introduction": "What is the circumnuclear CO morphology and how does it relate to the properties of the barred potential? The molecular gas shows a wide variety of  morphologies (Fig.~2) ranging from   relatively axisymmetric annuli or disks (starbursts NGC\\,4102, NGC\\,3504, NGC\\,4536,   and non-starbursts NGC\\,4314), elongated double-peaked  and spiral morphologies (starburst NGC\\,2782 and non-starbursts NGC\\,3351 and NGC\\,6951)  to extended distributions elongated along the large-scale bar (non-starburst NGC\\,4569). In  NGC\\,4569, the gas extends out to a large (2\\,kpc) radius, at a similar P.A. as the large-scale stellar bar (Fig.~2), and shows complex   non-circular motions. The   optical, NIR, and CO properties of NGC\\,4569 suggest it  is in the early stages of bar-driven/tidally-driven inflow of gas towards the inner kpc. In the other  galaxies, the gas distribution is less extended,  and in many systems it is concentrated inside the outer inner Lindblad resonance (ILR)  of the large-scale bar. As shown in Fig.~3, both starbursts and non-starbursts host ILRs. In the sample, the bar pattern speed $\\Omega_{\\mathrm{p}}> 40$--115\\,km~s$^{-1}$~kpc$^{-1}$, the  radius of the outer ILR is typically $>$\\,500\\,pc, and the radius of the inner IILR  $<$\\,300\\,pc. Note that  in NGC\\,2782 and NGC\\,470 which are claimed to host nuclear stellar bars \\cite{JKS99,FWRB96}, there is  a strong misalignment ($\\ge$\\,40\\,$^\\circ$) between the CO distribution and both the major axis and  minor axis of the large-scale stellar bar/oval.  \\begin{figure}[h] \\begin{center} \\includegraphics[width=2.0in]{jogee_fig3a.eps} \\includegraphics[width=2.0in]{jogee_fig3b.eps} \\end{center} \\caption[]{ ($\\Omega - \\kappa/2$) is plotted against radius. The bar pattern speed $\\Omega_{\\mathrm{p}}$  is drawn as horizontal  lines and estimated  by assuming that the corotation resonance is  near the end of the bar. Under the epicycle theory for a weak bar, the intersection of ($\\Omega - \\kappa/2$) with $\\Omega_{\\mathrm{p}}$ defines the  locations of the ILRs. } \\end{figure} ",
        "conclusions": ""
    },
    "0201/astro-ph0201491_arXiv.txt": {
        "abstract": "The recent results from a deep Fabry-Perot survey of nearby active and star-forming galaxies are presented.  Line-emitting material is detected over two orders of magnitude in galactocentric radius, from the 100-pc scale of the active or starbursting nucleus out to several 10s of kpc, sometimes well beyond the optical confines of the host galaxies.  The excitation and dynamical properties of the nuclear gas are studied to constrain the impact of galactic winds on the host galaxies and their environment. The properties of the warm ionized material on the outskirts of galaxies provide important clues for understanding galaxy formation and evolution. A new technique to search for starburst-driven wind galaxy candidates is discussed. The next generation of Fabry-Perot instruments on large telescopes promises to improve the sensitivity of emission-line galaxy surveys at least tenfold. ",
        "introduction": "The Fabry-Perot interferometer has the distinct advantage of providing detailed spectrophotometric information over a larger field of view (FOV) than that of other 3D instruments. The Fabry-Perot interferometer is therefore ideally suited to study nearby galaxies where the line-emitting gas extends over several arcminutes. This ionized material is an excellent probe of the phenomena taking place in the core of starburst and active galaxies, and can be used to quantify the impact of nuclear and star-formation activity on the environment and vice-versa.  Over the last decade, our group has used the scanning mode of the Hawaii Imaging Fabry-Perot Interferometer (HIFI) on the CFHT 3.6m and UH 2.2m and the TAURUS-II system on the AAT 3.9m and WHT 4.2m to study in detail a sample of about twenty nearby starbursts and active galaxies. A summary of the results from this portion of the survey is presented in \\S 2.  More recently, we have used a low-order Fabry-Perot interferometer to obtain very deep emission-line maps of several normal and active galaxies. This tunable-filter mode is particularly efficient to search for warm ionized gas on the outskirts of galaxies. The first results from this work are discussed in \\S 3. The tunable filter has also proven to be an efficient tool to search for starburst galaxies with large-scale galactic winds. This technique was used recently to uncover a shock-excited wind in the starburst galaxy NGC~1482. The data are described in \\S 4 along with possible applications of this technique to search for wind candidates at higher redshifts. Future avenues of research with Fabry-Perot interferometers are briefly discussed in \\S 5. ",
        "conclusions": ""
    },
    "0201/astro-ph0201172_arXiv.txt": {
        "abstract": "Deep, ground based, optical wide-field supernova searches are capable of detecting a large number of supernovae over a broad redshift range up to $z~\\sim~1.5$. While it is practically unfeasible to obtain spectroscopic redshifts of all the supernova candidates right after the discovery, we show that the magnitudes and colors of the host galaxies, as well as the supernovae, can be used to select high-$z$ supernova candidates, for subsequent spectroscopic and photometric follow-up.  Using Monte-Carlo simulations we construct criteria for selecting galaxies in well-defined redshift bands. For example, with a selection criteria using $B-R$ and $R-I$ colors we are able to pick out potential host galaxies for which $z~\\ge~0.85$ with 80\\% confidence level and with a selection efficiency of 64--86\\%. The method was successfully tested using real observations from the HDF.  Similarly, we show that that the magnitude and colors of the supernova discovery data can be used to constrain the redshift. With a set of cuts based on $V-R$ and $R-I$ in a search to $m_I~\\sim~25$, supernovae at $z \\sim 1$ can be selected in a redshift interval  $\\sigma_z \\le 0.15$. ",
        "introduction": "High-$z$ Type Ia supernovae (hereafter SNe) have been shown to be very accurate tools for studying cosmological parameters \\citep{per99,rie98}. As the sensitivity of the magnitude-redshift relation increases with the size of the redshift range probed, searches for SNe at redshifts $z~\\gsim$~1 become crucial to refine our understanding of cosmology \\citep{goo95,gol01}.  One of the difficulties in searches dedicated to high-$z$ SNe is to select candidates for subsequent follow-up. An $I$-band SN search down to $m_I~\\sim~25$ generates candidates in a broad redshift range: $0.1~<~z~\\lsim~1.5$. On a 1 square degree detection image, separated by three weeks from a reference image, there will be about a dozen Type Ia SNe on the rising part of the light curve.  Additional information is required to make the follow-up of $z~>$ 0.85 SNe efficient. This is particularly important as the follow-up of candidates with space instruments, such as HST, needs to be decided upon very soon after discovery, leaving limited time for prior spectroscopic screening.  In this paper we discuss how broad-band photometry of the host galaxy, as well as the SN can be used in order to select targets with $\\sigma_z \\sim 0.15$ with high efficiency. We focus on the optical bands, producing mainly SNe up to $z~\\sim~1.5$. In Section 2 we discuss the SN rates that are used in our simulations. In Section 3 we study the use of host galaxy magnitudes and colors to estimate the redshift of a SN candidate, followed by a discussion in Section 4 on how deep a survey must be in order to obtain the required galaxy colors. In section 5 we discuss the use of SN magnitude and color for redshift estimation. Finally, a summary and conclusions are given in Section 6.  In this paper we assume a flat cosmology with $\\Omega_M$ = 0.3 and $\\Omega_{\\Lambda}$ = 0.7, and a Hubble constant $H_0 = 65$ km s$^{-1}$ Mpc$^{-1}$ Magnitudes are given in the standard Vega based system. ",
        "conclusions": "We have shown that deep wide-field SN searches will detect a number of SNe over a broad redshift range. E.g., an optical search to $m_I~\\sim~25$ will find Type Ia SNe to $z~\\sim~1.5$. In this paper we discuss the use of magnitudes and colors of the SNe, as well as the host galaxies, in order for a prompt selection of high-$z$ SN candidates, for subsequent spectroscopic follow-up. We first discuss the use of additional photometric information on the host galaxies. The advantage of using the host galaxies for redshift estimates is that a catalog of high-$z$ candidate galaxies can be constructed in advance, enabling a direct selection once a SN is discovered.  Using the available magnitude information in the search filter gives a first indication of the probability that a SN is at high $z$. Adding broad-band filters increases the probability for a correct selection. Due to the cosmological shift of galaxy spectra, galaxies appear redder when they are at higher redshift (in optical to $z~\\sim~1$). Using the color from two filters, together with the magnitude information, one can therefore set additional constraints on the redshift of the galaxy. Note, however, that different galaxy types have different colors, as well as evolution in color, which makes a single color difficult to use. The selection efficiency increases if the galaxy type is known, however, one must remember that there is no one-to-one relation between morphological type and colors.  If observations of the host galaxies are available in three filters, then it is possible to derive selection criteria from the position of the galaxy in a color-color diagram. As an example, we show how such criteria can be constructed to select host galaxies with $z~>~0.85$. Using $B$, $R$ and $I$ photometry, we construct a selection criteria so that the probability is 80\\% that a galaxy has $z~>~0.85$, and where $\\sim~64-86$ \\% of the total number of host galaxies above this redshift are actually selected. By changing the selection criteria it is possible to increase the probability that the host galaxy is at high-$z$, this will, however, lead to a decrease of the fraction of the total number of high-$z$ host galaxies that are selected.  With four or more available filters it is possible to estimate photometric redshift of the host galaxies. This results in an estimate of the redshift for each galaxy separately, as compared to the three filter case, where we instead calculate the probability that a host galaxy has a redshift above, or below, some chosen limit.  We also discuss the necessary depth of a search in order to detect host galaxies at $z~\\sim~1$, and the depths required in other filters in order to calculate colors. From the results listed in Table \\ref{Table1} \\& \\ref{Table3}, we see that with a limiting magnitude $m_I~=~25$, about 80\\% the host galaxies at $z~\\sim~1.0$ can be detected, while at $z~\\sim~1.3$ the fraction is $\\sim$ 60\\%. The required depths in the other optical filters to detect 90\\% of the galaxies with $m_I~=~25$ are $m_B~\\sim~28.1$, $m_V~\\sim~26.8$ and $m_R~\\sim~25.6$.  In a similar way, we show that colors of the Type Ia SNe can also be used to constrain the redshift. A selection criteria based on $(V)RI$ magnitudes in a SN search with limiting magnitude $m_I~\\sim~25$, leads to a possible narrowing of the redshift range of the candidates to $\\sigma_z \\le 0.15$ at $z~\\sim~1$. Here 30-40\\% of the total number of SNe are selected, of which 75-90\\% have $z~>~0.9$. Multiband photometry of SNe is also very useful to distinguish between SNe of different types.  Selections based on the host galaxy magnitude and colors are preferable as they are less likely to introduce selection biases. Also, obtaining colors of the SNe requires deep multi-band  observations of each candidate taken within a relatively short time interval, making this technique less feasible.  More observational data is needed in order to understand and quantify the possible color-brightness correlation of high-$z$ SNe and the resulting selection bias.  Finally, the accuracy of the redshift determination, both form the galaxy and SN photometric information, depends on several parameters: depths, number of available filters, accuracy in photometry etc, and is a trade-off between required accuracy and the amount of observing time available."
    },
    "0201/astro-ph0201458_arXiv.txt": {
        "abstract": "Parsec scale Faraday rotation measure maps are presented for the radio galaxies M87, 3C\\,111, and 3C\\,120. These VLBA observations were made at 8, 12, and 15 GHz. M87 has an extreme RM distribution which varies from $-$4000 \\radm\\ to more than 9000 \\radm\\ across a projected distance of 0.3 parsecs in its jet. M87 has no polarized flux closer than 17 mas from the core. 3C\\,111 and 3C\\,120 both show polarized emission in their cores which is consistent with the expectations of unified schemes for these broad line radio galaxies. 3C\\,111 has an RM gradient which increases from $\\sim -200$ \\radm\\ 4 mas from the core to $\\sim -750$ \\radm\\ on the side of the jet closest to the core. 3C\\,120 has a more moderate RM distribution in the jet of approximately 100 \\radm\\ but this increases by an order of magnitude in the core. ",
        "introduction": "Several recent papers (Udomprasert et al. 1997, Cotton 1997 and Taylor 1998 \\& 2000) have demonstrated that extreme values of Faraday rotation of up to 40000 \\radm\\  in the rest frame of quasars are possible. \\citet{zt01} showed that the Rotation Measure (RM) properties of quasars vary on both small spatial (parsec) scales and short timescales (1.5$-$3 years). These observations suggest that the observed RM distributions are intrinsic to the central few hundred parsecs of AGN and are not a foreground effect from the host galaxy ISM, the ICM, or the Milky Way. Hence, the observed Faraday rotation can serve as a probe of the magnetic field weighted by the electron density along the line of sight near the central engines of AGN. With the electron density supplied by spectral line diagnostics (e.g. Osterbrock 1989) an estimate of the magnetic field strength and orientation can be made provided one assumes a physically reasonable path length.  To date parsec-scale RM observations have only been made for quasars. Polarimetric observations of other classes of AGN are required in order to test the unified model for AGN \\citep{ant93} by looking for orientation effects in the RM distributions. In order to form a more complete picture of the parsec scale RM properties of AGN \\citet{tay00} designed a sample of 40 quasars, BL Lacs, and radio galaxies for polarimetric observations. We have completed this survey with the VLBA and present here the first sub-parsec scale RM maps for radio galaxies.  We assume H$_0 = 50$ km s$^{-1}$ Mpc$^{-1}$ and q$_0$=0.5 throughout. ",
        "conclusions": ""
    },
    "0201/astro-ph0201344_arXiv.txt": {
        "abstract": "An exquisite gravitational arc with a radius of 2\\farcs1 has been discovered around the $z = 0.938$ field elliptical galaxy CFRS03.1077 during $HST$ observations of Canada-France Redshift Survey (CFRS) fields.  Spectroscopic observations of the arc show that the redshift of the resolved lensed galaxy is z = 2.941.  This gravitational lens-source system is well-fitted using the position angle and ellipticity derived from the visible matter distribution and an isothermal mass profile with a mass corresponding to $\\sigma=387\\pm5$ km s$^{-1}$. Surprisingly, given the evidence for passive evolution of elliptical galaxies, this is in good agreement with an estimate based on the fundamental plane for $z=0$ ellipticals. This, perhaps, indicates that this galaxy has not shared in the significant evolution observed for average elliptical galaxies at z$\\sim$ 1. A second elliptical galaxy with similar luminosity from the CFRS survey, CFRS 14.1311 at $z=0.807$, is also a lens but in this case the lens model gives a much smaller mass-to-light ratio, i.e., it appears to confirm the expected evolution.  This suggests that this pair of field elliptical galaxies may have very different evolutionary histories, a significant result if confirmed.  Clearly, CFRS03.1077 demonstrates that these ``Einstein rings\" are powerful probes of high redshift galaxies. ",
        "introduction": "A very beautiful arc structure, characteristic of a lensed object, was noted around the galaxy CFRS03.1077 \\citep{Ham95} during analysis of $HST$ images of high redshift field galaxies. CFRS03.1077 is a bright elliptical galaxy at z = 0.938 and hence is likely to lens background galaxies or quasars. \\citet{Mir92} predicted that $\\sim$100 gravitationally lensed rings per square degree should be observable at optical wavelengths to $B \\sim 26$, given the number of Einstein rings at radio wavelengths and the projected density of galaxies in optical surveys.  Indeed, \\citet{Rat99} report finding 10 good galaxy lens candidates in 0.1 square degrees, based on analyses of $HST$ Medium Deep Survey images. These, however, have not yet been confirmed and remain only as candidates.  It is now recognized that galaxies which act as gravitational lenses are extremely important tools for studying a variety of cosmological problems including the nature and evolution of the lensing galaxies themselves \\citep{Koc00,Nar99,Ref94, Schn96}. For example, \\citet{Im97} showed that observations of seven galaxy lenses  favor a nonzero cosmological constant, while \\citet{Koc00} and \\citet{Kee98} present results on the properties of the lensing galaxies, most of which are early-type. Since the lens models give data on the total (dark + visible) distribution of mass, investigations like the latter are extremely powerful, especially for studying the properties of galaxies at high redshifts where traditional methods become increasingly difficult. Einstein ring lenses provide additional constraints and hence are even more important to the determination of mass distributions of the lensing galaxies \\citep{Koc01}.  Even though Einstein ring lenses are predicted to be much more common at optical wavelengths than radio, only one optically-identified system has so far been confirmed:  0047-2808 \\citep{War98, War99}. Rather than being initially recognised from an image as a ``ring\", this lens was detected through superposition of emission lines from the lensed object on the spectrum of the foreground galaxy. \\citet{Hew00} and \\citet{Hal00} recently report finding additional candidate lenses in this manner, but these remain to be confirmed. These authors argue that spectroscopic observation of distant early type galaxies is one of the best methods of detecting galaxy lenses, particularly because the presence of spectroscopic features from the source implies that the crucial redshifts of both the source and deflector can be determined. In practice, however, this method is only viable for background galaxies with strong emission lines.  Furthermore, galaxy lens candidates are quite easily recognizable with images from $HST$ or ground-based adaptive optic systems, and disentangling the spectral features arising from the source and deflector galaxies is no longer as difficult with the improved spatial resolution being achieved on modern large ground based telescopes.  In this paper we report on a detailed analysis of the $HST$ images of the lens and source galaxies related to CFRS03.1077. We also present spectra of these galaxies obtained with the Canada-France-Hawaii Telescope (CFHT) and derive the redshift of the source - a galaxy without strong emission lines. Finally, a simple model of the lens-source system is described. ",
        "conclusions": "Our spectroscopic observations confirm that the arc surrounding the z = 0.938 elliptical galaxy CFRS03.1077 is indeed a lensed image of a background galaxy. The redshift of this galaxy is z = 2.941. Standard lens models easily reproduce the observed arc structure and also suggest that two faint objects observed near the lensing galaxy on the opposite side to the arc are lensed images. Observations at other wavelengths should be obtained to determine the colors of these objects. If they are the same as the arc, then this would be further evidence that they are lensed images and can be confidently used to constrain the lens geometry. Multi-wavelength observations of CFRS03.1077 could also be used to examine whether the internal colors of the galaxy itself are normal or show strong variations indicative of recent star formation (cf \\citet{Men01}). If the internal colors are not homogeneous this may help explain why the line of sight velocity dispersion determined from the lens model is higher than expected from fundamental plane considerations assuming passive evolution since $z \\sim 1$.  CFRS03.1077 demonstrates the potential offered by detailed study of Einstein ring lenses. With the advent of 8m class telescopes, especially those equipped with integral field spectrographs, the spectroscopic data reported here can now be very significantly improved. Images at other wavelengths should be obtained to establish or identify additional lensed images in order to more tightly constrain the lens model."
    },
    "0201/hep-ph0201209_arXiv.txt": {
        "abstract": "In low scale quantum gravity scenarios the fundamental scale of nature can be as low as TeV, in order to address the naturalness of the electroweak scale. A number of difficulties arise in constructing specific models; stabilisation of the radius of the extra dimensions, avoidance of overproduction of Kaluza Klein modes, achieving successful baryogenesis and production of a close to scale-invariant spectrum of density perturbations with the correct amplitude. We examine in detail the dynamics, including radion stabilisation, of a hybrid inflation model that has been proposed in order to address these difficulties, where the inflaton is a gauge singlet residing in the bulk. We find that for a low fundamental scale the phase transition, which in standard four dimensional hybrid models usually ends inflation, is slow and there is second phase of inflation lasting for a large number of $e$-foldings. The density perturbations on cosmologically interesting scales exit the Hubble radius during this second phase of inflation, and we find that their amplitude is far smaller than is required.  We find that the duration of the second phase of inflation can be short, so that cosmologically interesting scales exit the Hubble radius prior to the phase transition, and the density perturbations have the correct amplitude, only if the fundamental scale takes an intermediate value.  Finally we comment briefly on the implications of an intermediate fundamental scale for the production of primordial black holes and baryogenesis. ",
        "introduction": "In nature there are two apparent scales; the electroweak scale and the scale of gravity, separated by seventeen orders of magnitude. Understanding the gap between these scales has been a prime motivation behind studying theories beyond the electroweak Standard Model (SM). Supersymmetry provides an elegant scheme for keeping the electroweak scale stable under any large radiative corrections, however the lack of direct evidence for supersymmetry in collider physics and in nature has lead to the consideration of scenarios with large extra dimensions. In these scenarios the fundamental scale is taken to be the higher dimensional Planck mass, $M_{\\ast}$, which is assumed to be close to the electroweak scale~\\cite{nima0,early}. While in this scheme supersymmetry is redundant in four dimensions, the presence of low energy supersymmetry could still be a viable option, however, with the fundamental scale at an intermediate scale, somewhere between the Planck and electroweak scales. Such a scenario is well motivated by string theory~\\cite{quevedo}, which predicts that gauge and gravity unification occurs below the Grand Unification Scale $\\sim 10^{16.5}$GeV.   The four dimensional Planck mass in these theories is obtained via dimensional reduction, assuming that the extra dimensions are compactified on a torus, the simplest possible manifold. The volume of the extra dimensions $V_d$, the effective four dimensional Planck mass and the fundamental scale are then simply related: \\begin{equation} \\label{imp} M_{\\rm p}^2 = M_*^{2+d} V_d \\,, \\end{equation} where $d$ is the number of extra compact dimensions. For given $M_{\\ast}$ this fixes the present day size of each of the extra dimensions, $b_0$. For two extra dimensions and $M_{\\ast} \\sim 1$ TeV, $b_0 \\sim 0.2$ mm. Currently collider physics and supernovae $1987A$ impose a bound on the fundamental scale: $M_{\\ast}\\geq 30$ TeV \\cite{nima0,exp1,exp2}, and a recent astrophysical bound based on a heating of a neutron star suggets $M_{\\ast} \\geq 500$~TeV \\cite{new}. If the fundamental scale is as low as $\\sim 100$ TeV, it is important that the SM particles are trapped in a four dimensional hypersurface (a $3$-brane) and are not allowed to propagate in the bulk~\\cite{nima0}. It is generically assumed that besides gravity, the SM singlets, which may include the inflaton, can propagate in the bulk \\cite{abdel100}.   The cosmological setup in models with large extra dimensions is quite different from the conventional one. Firstly, if the electroweak scale is the fundamental scale in higher dimensions, then there can be no massive fields beyond the electroweak scale in four dimensions. Secondly, the size of the extra dimensions can be quite large, compared to the electroweak scale, which implies the existence of new degrees of freedom, usually known as the radion, with a mass scale as small as ${\\cal O}(0.01 {\\rm eV})$ if there are two large extra dimensions. The large extra dimensions must grow from their natural scale of compactification, $\\sim(\\rm TeV)^{-1}$, and then stabilize at around a millimeter. This stabilization must occur before the electroweak phase transition and nucleosynthesis, via some kind of a trapping mechanism as discussed in Ref.~\\cite{abdel2}. The Kaluza Klein (KK) states of the graviton, and any other fields residing in the bulk, can be excited at high temperatures and hence lead to constraints on these models. Above the {\\em normalcy temperature}, the Universe could be filled by the KK modes. For Big Bang Nucleosynthesis to occur successfully the normalcy temperature must be greater than $\\sim 1$ MeV. Furthermore the final reheat temperature, which is constrained by cosmological considerations to be as small as $100$ MeV~\\cite{nima0,davidson0,aemp}, should be smaller than the normalcy temperature. These considerations severely restrict baryogenesis in these models, for a detailed discussion see Refs.~\\cite{abdel3,aemp}.   Constructing a successful inflation model, which produces a close to scale invariant spectrum of density perturbations with the correct amplitude and a very low reheat temperature is a challenging issue, with single field models and models where the inflaton is a brane field proving particularly problematic~\\cite{many,anu}. There have been several proposals~\\cite{many,anu}, arguably the most natural of which invokes SM singlet scalars coupled together to form a potential which mimics that of the standard four dimensional hybrid inflationary model, but with the fields promoted to the higher dimensions~\\cite{abdel1,abdel2}. It has also been shown that baryogenesis can occur successfully in this model~\\cite{abdel3,aemp}. In this paper we study the dynamics of this extra dimension inspired hybrid inflationary model in detail.    In hybrid inflation models, the false vacuum field is initially trapped in a stable minimum at zero whilst the inflaton field slow-rolls down its potential. At some critical value of the inflaton field the stable minimum becomes an unstable maximum and quantum diffusion produces a second order phase transition from the false vacuum to the true vacuum~\\cite{linde}. In standard four dimensional cosmology, for most parameter values, the bare mass of the false vacuum field is much greater than the Hubble parameter and the phase transition occurs rapidly and inflation ends. If these quantities have roughly the same magnitude, however, then the roll-down of the false vacuum field is no longer fast and a second period of inflation occurs~\\cite{randall,garcia}. In standard four dimensional cosmology, for the phase transition to occur slowly the effective coupling of the false vacuum field has to be tiny, $\\sim 10^{-30}$~\\cite{garcia}. We find, however, that for the parameter values which are relevant for the extra dimensional model (fundamental coupling constants of order unity and fundamental scale $\\sim 100$TeV~\\cite{abdel1,abdel2}), the bare mass of the false vacuum field is of order the Hubble parameter so that the phase transition is slow and a second period of inflation occurs.   The duration of this second phase of inflation is typically very long, so that the density perturbations on cosmological scales are generated close its end. We calculate the amplitude of these perturbations and find that they are only compatible with the COBE normalization if the fundamental scale is significantly larger than 100 TeV. Finally we discuss the implications, specifically achieving successfully baryogenesis and avoiding the over-production of primordial black holes (PBHs), of an intermediate fundamental scale. In order to keep our discussion as general as possible, we do not fix either the fundamental scale or the number of extra dimensions from the outset, however we will focus throughout on the parameter values of the specific model proposed in Refs.~\\cite{abdel1,abdel2}. ",
        "conclusions": "We have studied in detail the dynamics of a hybrid inflationary model in the context of large extra dimensions, proposed in Ref.~\\cite{abdel1} and subsequently studied in Refs.~\\cite{abdel2,abdel3,aemp}, where it is assumed that the inflaton sector is a gauge singlet residing in the higher dimensions. We have studied the entire gamut of the dynamical behavior of the coupled scalar fields. In particular we have shown that for a low fundamental scale, as studied in Refs.~\\cite{abdel1,abdel2,abdel3}, there are in fact two distinct phases of inflation.  The first phase of inflation has two parts, an initial period of radion dominated inflation, as the radion is stabilised via dynamical trapping in its own potential, followed by vacuum dominated inflation, as in standard hybrid models. At low energies, when the radion mass dominates the Hubble parameter, the radion begins oscillating around the minimum of its potential, however if the bare mass of the radion is very small, $\\lesssim 10^{-2}$eV in the case of two large extra dimensions and the fundamental scale ${\\cal O} ({\\rm TeV})$, then the radion density stored in the oscillations is not large enough to cause problems similar to the moduli problem.   For the parameter values considered natural in standard four dimensional hybrid inflation models, inflation ends rapidly, once the $\\phi$ field reaches the critical value at which the false vacuum becomes unstable.  For the parameters which are natural from an extra dimensional perspective (a fundamental scale $M_{\\ast} \\sim N_0 \\sim 10^{5}$ GeV and fundamental coupling constants $\\lambda,g \\sim {\\cal O}(1)$) we find that the phase transition is slow, producing a second phase of inflation which lasts for around $10^{6}$ $e$-foldings. Cosmologically interesting scales therefore exit the Hubble radius close to the end of this second period of inflation, and we find that the amplitude of the density perturbations on these scales is smaller than required by the COBE normalisation unless the vacuum expectation value of the false vacuum field, $N_{0}$, is less than $10^{-4}$ GeV.  The only way around this obstacle is to try to shorten the second phase of inflation so that cosmologically interesting scales exit the Hubble radius before the phase transition. To do this we need a fundamental scale higher than $10^{5}$GeV. We have found that, for instance, with the set of parameter values $M_{\\ast}\\sim 10^9$GeV, $N_{0}\\sim 10^4$GeV and $m_{\\phi}\\sim 10^{-5}$GeV density perturbations of the correct amplitude are produced. For this intermediate fundamental scale the constraint on the normalcy temperature is relaxed so that electroweak baryogenesis and leptogenesis become possible, furthermore the baryogenesis mechanism provided in Ref.~\\cite{aemp} is unaffected by the second phase of inflation and remains viable. If there is a second phase of inflation lasting less than 43 $e$-foldings then the large density perturbations produced at the beginning of the phase transition will have re-entered the Hubble radius by the present day, and may lead to the over-production of PBHs~\\cite{randall,garcia}. The formation criteria for PBHs in extra-dimensional scenarios have not yet been studied, so it is not possible to use the constraints on PBH abundance~\\cite{pbh} to constrain the model parameters."
    },
    "0201/astro-ph0201350_arXiv.txt": {
        "abstract": "The FREGATE experiment aboard HETE-II has been successfully integrated into the Third Interplanetary Network (IPN) of gamma-ray burst detectors.  We show how HETE's timing has been verified in flight, and discuss what HETE can do for the IPN and vice-versa. ",
        "introduction": "The FREGATE experiment aboard HETE-II is an excellent complement to the detectors in the IPN.  It has good sensitivity to bursts, thanks to its large surface area and its steady background in equatorial orbit.  Its time resolution in both triggered and untriggered modes is sufficiently high for cross-correlating precisely with other spacecraft.  And finally, its energy range is well matched to those of other instruments in the network.  However, before any experiment can be added to the network, its timing must be verified in flight.  In this paper, we will first explain the procedure we have used to demonstrate that the HETE timing is accurate.  Then we will show how the IPN can be used to improve the location accuracy of HETE bursts, and how the FREGATE data are useful to the IPN. ",
        "conclusions": "We have successfully integrated the HETE spacecraft into the third interplanetary network of gamma-ray burst detectors, and have demonstrated the value of HETE to the IPN and vice-versa.  We are now looking forward to several more years of operations.   KH is grateful for IPN support under JPL Contract 958056, and for HETE-FREGATE support under MIT Contract SC-R-293291."
    },
    "0201/astro-ph0201399_arXiv.txt": {
        "abstract": "We report the discovery of the optical and radio afterglow of GRB~010921, the first gamma-ray burst afterglow to be found from a localization by the High Energy Transient Explorer (HETE) satellite. We present optical spectroscopy of the host galaxy which we find to be a dusty and apparently normal star-forming galaxy at $z = 0.451$.  The unusually steep optical spectral slope of the afterglow can be explained by heavy extinction, $A_V > 0.5$ mag,  along the line of sight to the GRB.  Dust with similar $A_V$ for the the host galaxy as a whole appears to be required by the measurement of a Balmer decrement in the spectrum of the host galaxy.  Thanks to the low redshift, continued observations of the afterglow will enable the strongest constraints, to date, on the existence of a possible underlying supernova. ",
        "introduction": "The High Energy Transient Explorer (HETE-2) was successfully launched and deployed on 2000 October 9 \\citep{ricker+00}.  As the first satellite entirely dedicated to the detection and study of gamma-ray bursts (GRBs), HETE's primary mission is the localization of GRBs through their prompt emission and the rapid relaying of the coordinates to ground-based observers.  Here, we present the discovery of the afterglow of GRB~010921 \\citep{ricker+01b}, the first afterglow to be identified through localization by HETE-2. ",
        "conclusions": "\\label{sec:conclusion}  Here we report the first afterglow resulting from an HETE localization of a GRB.  The afterglow in itself is similar to afterglows studied to date. However, the GRB occured in a low-redshift, $z=0.451$, dusty star-forming galaxy. Such low redshift GRBs are valuable for two reasons. First and foremost, such low-redshift events offer the best opportunity to study the explosion physics including the opportunity to constrain the presence of (or detect) any underlying supernova. Second, the low redshift allowed us to determine the metallicity of the host galaxy with relative ease.  It is perhaps significant to note that the afterglow was discovered through the use of image differencing.  We suggest that the use of this technique may be a more robust method of identifying GRB afterglows than the traditional manual comparison with sky survey images, since it enables the detection of an afterglow superimposed on a bright host galaxy. This is particularly important for low redshift GRBs.  We wonder whether the traditional approach --- looking for an isolated transient --- could be the cause of failure to identify optical afterglows in some GRBs."
    },
    "0201/astro-ph0201166_arXiv.txt": {
        "abstract": "The dense central region of NGC~6397 contains three classes of stars whose origins are likely related to stellar interactions: blue stragglers (BSs), cataclysmic variables (CVs), and probable helium white dwarfs (HeWDs).  We summarize results to date concerning CVs and HeWD candidates that have been identified in two imaging studies with Hubble Space Telescope (HST), and present one new CV candidate that appears well outside the cluster core.  We also present results concerning binaries containing two main sequence stars in the central parts of the cluster.  Proper motion information derived from two epochs of HST data is used to remove field stars from the sample. Binaries are then identified on the basis of their positions in the color-magnitude diagram.  We set an upper limit of \\about 3\\% on the fraction of main sequence stars with primary masses in the range $0.45-0.8$\\msun\\ and mass ratios q $\\ga$ 0.45.  Extrapolating to all mass ratios gives an estimated binary fraction of $\\la$ $5-7$\\%.  Even in these small numbers, such pairs are likely to be key players in the processes that give rise to the more exotic stellar populations. ",
        "introduction": "The nearest places in the Universe where stellar collisions should be relatively commonplace are in the central regions of globular clusters with collapsed cores.  The nearest such cluster is NGC~6397.  At a distance of \\about $2.7\\pm 0.2$ kpc (Reid 1998, Reid \\& Gizis 1998, Anthony-Twarog \\& Twarog 2000), its dense central region is both brighter and more spread out on the sky than other high-density cluster cores.  NGC~6397 is thus a prime locus in which to study the consequences of stellar interactions.  The inner regions of NGC~6397 are characterized by a power-law surface-brightness profile surrounding a resolved core with a density of \\about 10$^5$\\msun\\ pc$^{-3}$.  Using star counts in HST/WFPC2 images, Sosin (1997) derived a maximum-likelihood core radius of 4.8\\asec (1\\asecspt 5$-8$\\asec\\ at 95\\% confidence) for stars with masses close to the turnoff mass, consistent with values determined from the ground (e.g., Drukier 1993).  Sosin's analysis also revealed a high degree of mass segregation in the central regions, including a mass function that in the inner \\about 18\\asec\\ is inverted, i.e., increases with increasing mass.  The first evidence for unusual stellar populations in NGC~6397 was the discovery of a central concentration of bright blue stragglers (BSs) by Auri\\` ere, Ortolani, \\& Lauzeral (1990).  Since then, high-resolution imaging with the Wide Field and Planetary Camera 2 (WFPC2) on Hubble Space Telescope (HST) has revealed two more classes of stars whose origins are also likely to be linked to stellar interactions: cataclysmic variables (CVs) and probable helium white dwarfs (HeWDs).  The presence of CVs was first signaled by a population of faint X-ray sources detected with ROSAT (Cool \\et\\ 1993).  Several optical counterparts were subsequently identified via \\ha\\ excess, UV excess, and/or variability (De Marchi \\& Paresce 1994, Cool \\et\\ 1995, 1998), and confirmed spectroscopically (Grindlay \\et\\ 1995).  More recently, a population of blue stars similar in brightness to CVs was found that lacked the photometric variability of CVs (Cool \\et\\ 1998).  Dubbed ``non-flickerers'' (NFs), they have been identified as probable helium white dwarfs (HeWDs) (Edmonds \\et\\ 1999).  \\begin{figure} \\plotfiddle{cool_bolton_fig1.ps}{4.0in}{0}{65}{65}{-210}{-130} \\caption{CMDs of stars within 57\\asec\\ of the cluster center (left) vs.\\ stars outside 57\\asec\\ (right).  Equal numbers of stars appear in these two regions of the WFPC2 field.  Solid dots mark the five bright blue stragglers in the cluster core.  The bluest of these is the one that lies closest to the cluster center and was found by Saffer \\et\\ (this volume) to be most massive.  X's mark the fainter blue stragglers, all of which lie at larger radii.  Triangles and squares mark cataclysmic variables and probable helium white dwarfs, respectively.} \\end{figure}  \\begin{figure} \\plotfiddle{cool_bolton_fig2.ps}{4.5in}{0}{75}{75}{-230}{-180} \\caption{Locations within the WFPC2 field of cataclysmic variables, helium white dwarfs candidates, and blue stragglers.  Symbols are as in Fig.~1.  The massive blue straggler (Saffer \\et, this volume) is the one closest to the center of the cluster (middle one of the five large dots).  The center and core radius adopted are those of Sosin (1997): R.A. $= 17^{\\rm h} 40^{\\rm m} $42\\secspt 06, Dec. $= -53$\\deg\\ 40\\amin\\ 28\\asecspt 8 (uncertainty \\about 1.5\\asec), and r$_c$ = 4\\asecspt 8 (see text).  This center corresponds to (x,y) $=$ (529,378) on HST archive image u33r010kt (PC1 chip).  The open circle marks a star that Saffer \\et\\ identify as a probable horizontal-branch star; it had previously been identified as a blue straggler (Lauzeral \\et\\ 1992).} \\end{figure}  Here we present results from two epochs of HST/WFPC2 imaging of the central regions of NGC~6397.  Several of the results concerning CVs and NFs have appeared elsewhere.  We collect those results and present new color-magnitude diagrams (CMDs) showing all the CVs and NFs along with the BSs that appear in the field.  We also describe new findings regarding a sixth CV candidate, the first to be found many core radii from the cluster center.  We then discuss initial results of a study of main-sequence binaries in the cluster.  This study makes use of the two epochs of WFPC2 data to perform a proper-motion selection to remove field stars from the sample of main-sequence binary candidates.  With their large cross sections, such binaries are likely to play a critical role in stellar interactions, affecting the dynamical evolution of the cluster as well as the formation of exotic stellar populations.  Understanding the populations of BSs, CVs, and NFs is likely to require an understanding of main-sequence binaries as well.  Fig.~1 shows CMDs for all stars in the two epochs of WFPC2 data brighter than V = 22.5.  The BSs, CVs, and NFs are all marked.  We have divided the stars into two equally sized groups, one with r $<$ 57\\asec (left panel) and one with r $>$ 57\\asec\\ (right panel).  It can be seen at a glance that all three populations of stars are strongly concentrated toward the center of the cluster.  This central concentration can also be seen in Fig.~2, in which we plot the locations of the BSs, CVs, and NFs within the HST/WFPC2 field. ",
        "conclusions": "Stars in the middle of dense clusters must collide.  Newton's laws and the laws of probability require it.  But the outcomes of the collisions, and especially of the more numerous near misses, are less certain.  With HST it is possible to observe the products of these interactions directly.  In NGC~6397, every type of star seems to be in on the action.  Main sequence stars are combining to form blue stragglers; white dwarfs are coupling with main-sequence stars to form cataclysmic variables; giants are being stripped of their envelopes to reveal their helium cores.  The details of the particular processes that are generating this stellar mena\\-gerie will take time to sort out.  Multiple pathways exist by which each of these classes of stars may form, and the presence of binaries, even in the small numbers found here, multiplies the possibilities.  Formation rates, lifetimes, and destruction/ejection rates all need to be factored in to make sense of the variety and distributions of objects being found.  The growing sophistication of simulations of clusters and stellar interactions, coupled with improved constraints on the numbers and distributions of both ordinary and extraordinary stars in clusters, holds out the promise of significant progress in the near future.  New dynamical modeling of NGC~6397, making full use of the new information HST has provided, would clearly be valuable.  Here we draw attention to a few points of particular interest that may bear further investigation.  \\vskip 0.05in  (1) Why are there so few main-sequence binaries in NGC~6397?  \\vskip 0.05in  The limit that we place on the main-sequence binary fraction in NGC~6397 is one of the lowest reported for a globular cluster.  Most values, obtained using a variety of different techniques, range upward of 10\\% (see Hut \\et\\ 1992 for a review).  Particularly notable is the much higher fraction (15\\%--38\\%) measured by Rubenstein \\& Bailyn (1997) in the center of NGC~6752, another nearby, core-collapsed cluster with otherwise similar properties.  Perhaps NGC~6397 is farther along in the process of burning up its original store of binaries than NGC~6752.  The short relaxation time in NGC~6397 is likely to facilitate the rapid feeding of binaries into the core where they will quickly be modified through interactions with other stars (McMillan \\& Hut 1994, and references therein), perhaps even to form some of the blue stragglers now seen in the core.  But the present dynamical properties of NGC~6752 are not so different from those of NGC~6397 as to make this a particularly satisfying hypothesis.  Whether their past histories (e.g., the timing of core collapse) might be sufficiently different one can only speculate about, and the possibility remains that the two clusters were simply born with significantly different binary populations.  Still, it appears likely that the main-sequence binaries we now observe in NGC~6397 are a vestige of the original population.  Large interaction cross sections coupled with the short relaxation times in NGC~6397 (\\about 10$^5$ yr in the center; \\about $2\\times 10^8$ yr at the half-mass radius of 2\\aminspt 8---Djorgovski 1993) make binaries extremely vulnerable to destruction through a variety of direct and indirect processes.  Even if the binaries we are now seeing were formed when the cluster was born, it appears that few will been left unscathed. Judging from the calculations of Davies (1995), the rates of binary hardening and exchange interactions should be sufficiently high in this cluster that most systems will have significantly reduced periods, and some fraction of their present components will not be original.  Thus it remains to be seen, for example, whether the distribution of mass ratios that we derive has more to say about the cluster's primordial population of binaries or about the nature of exchange interactions.  Either way, the fraction of these binaries, while small, is large enough to contribute significantly to the production of CVs, NFs, and/or BSs (Davies 1995).  \\vskip 0.05in  (2)  What accounts for the distribution of BSs, CVs, and NFs?  \\vskip 0.05in  Saffer \\et\\ (this volume) have measured masses for the five bright blue stragglers in the core of NGC~6397 (see Fig.~2).  Four have masses of about twice the turnoff mass; the fifth is of order three times the turnoff mass.  Interestingly, the latter is the blue straggler closest to the cluster center---exactly where one might expect the most massive objects to reside (see Fig.~2).  Its position is consistent with being at the cluster center, according to the star count analysis of Sosin (1997) as well as the ground-based analysis of Auri\\` ere \\et\\ (1990).  The distribution of the remaining bright BSs is more puzzling.  Why are they so much more concentrated toward the center than the CVs, which ought to have comparable (though somewhat lower) masses?  The NFs present a puzzle of their own.  Within the limits of small-number statistics, their distribution appears similar to that of the CVs. Yet if they are He WDs, their masses are much too low (\\about 0.25\\msun) to account for this similarity.  A natural solution is that they have binary companions.  Companions are needed in any case to explain the formation of a He WD, whether through Roche-lobe overflow or via common-envelope evolution following the collision of a red giant with another star.  Sufficiently massive main-sequence companions are ruled out (Cool \\et\\ 1998), but neutron stars, or possibly massive white dwarfs, are viable.  Edmonds \\et\\ (1999) find preliminary evidence for binarity in the radial velocity of the one NF for which a spectrum has been obtained; additional studies are clearly needed.  This still leaves open the question of why the bright BSs are much more concentrated than either the CVs or NFs.  One factor that may distinguish the BSs from the CVs and NFs is binarity.  The CVs are certainly binaries; the NFs make little sense as He WDs unless they too are binaries.  The bright BSs show no evidence for photometric variability that would suggest binarity (in contrast to some of the fainter BSs in the cluster---Cool \\et\\ 2001).  Perhaps the recoil that binaries experience in interactions with other stars keeps them more spread out, on average, than single stars of comparable mass.  A further puzzle is that all three populations (BSs, CVs, and NFs) appear much more centrally concentrated than the population of 1.4\\msun\\ neutron stars that Dull (1996) included in his Fokker-Planck model of NGC~6397.  Yet their masses should all be roughly comparable. This suggests that not all of these objects have come into equilibrium with the other stars in the cluster, despite having lifetimes well in excess of the relevant 2-body relaxation times.  Meanwhile, Sosin's (1997) analysis of the epoch~1 HST data suggests that the main-sequence stars {\\it are} in energy equipartition with one another within the cusp (radius \\about 100\\asec).  It will be interesting to see whether this apparent discrepancy between the behavior of main-sequence stars and probable interaction products is borne out in future dynamical model that take account of the HST results.  \\vskip 0.05in  (3) Why is CV~6 out at \\about 15 core radii?  \\vskip 0.05in  Numerical simulations have shown that close binary stars in dense cluster cores are vulnerable to interactions with passing stars that can eject the binaries from the cluster.  If its recoil velocity is insufficient to escape from the cluster, the binary may spend considerable time at large radii before sinking back to the central region of the cluster.  If this has happened to CV~6, it appears that it would likely have been a relatively recent occurrence.  The central relaxation time in NGC~6397 is exceedingly short; even at the half-mass radius (\\about 170\\asec) it is only \\about 2\\x 10$^8$ yr. Unless its radial distance from the center is much greater than its 72\\asec\\ projected distance, CV~6 shouldn't take long to migrate back into the center.  In view of the possibility that CV~6 has been ejected from the center in the relatively recent past, it appears somewhat tantalizing at first glance that in the proper motion diagram of Fig.~5, CV~6 is the triangle offset from the rest of the CVs and NFs.  However, CV~6 is the only CV that lands on a WF chip rather than on the PC1 chip.  The larger WF pixels will result in a correspondingly larger uncertainty in the measurement of its proper motion.  Moreover, the direction of motion implied by the proper motion (assuming for the moment that the offset from (0,0) is significant) is closer to being tangential than radial.  We conclude that it is most likely that the offset of the proper motion of CV~6 from (0,0) is little more than measurement error.  Further refinements by Anderson and King of the astrometric techniques they have developed for WFPC2 may permit direct measurements of the internal motions of cluster stars in the near future (see, e.g., Anderson \\et\\ 1998).  Such measurements of this and the other CVs, NFs, and BSs in NGC~6397 would be of particular interest."
    },
    "0201/astro-ph0201485_arXiv.txt": {
        "abstract": "We have investigated the structural and dynamical properties of  triaxial stellar systems whose surface brightness profiles follow the $r^{1/n}$ luminosity law -- extending the analysis of Ciotti (1991) who explored the properties of spherical $r^{1/n}$ systems.  A new analytical expression   that accurately reproduces the  spatial (i.e.\\ deprojected)  luminosity density profiles (error less than 0.1\\%) is presented  for detailed modelling of the S\\'ersic family of luminosity profiles. We evaluate both the symmetric and the non--axisymmetric components of the gravitational potential and   force and compute the torques as a function of position. {\\it For a given triaxiality, stellar systems with  smaller values of $n$ have a  greater  non--axisymmetric gravitational field component}. We also explore the strength of the non--axisymmetric forces produced by bulges with differing $n$ and triaxiality on systems having a range of bulge--to--disc ratios. The increasing disc--to--bulge ratio with increasing galaxy type (decreasing $n$) is found to heavily reduce the amplitude of the non--axisymmetric terms, and therefore reduce the possibility that triaxial bulges in late--type systems may be the mechanism or perturbation for non--symmetric structures in the disc.    Using seeing--convolved $r^{1/n}$--bulge plus exponential--disc fits to the K--band data from a sample of 80 nearby disc galaxies, we probe the relations between galaxy type, S\\'ersic index $n$ and the bulge--to--disc luminosity ratio. These relations are shown  to be primarily a consequence of the relation between  $n$ and the total bulge luminosity.  In the K--band, the trend of decreasing bulge--to--disc luminosity ratio along  the spiral Hubble sequence is predominantly, although not entirely, a consequence of the change in the total  bulge luminosity;  the trend between the total disc  luminosity and Hubble type is much weaker. ",
        "introduction": "As the quality of photometric data has improved over the years (largely due to the use of CCDs), the applicability of a fitting-function which can account for  variations in the curvature of  a light profile has been demonstrated for elliptical galaxies  (Capaccioli 1987, 1989; Davies et al. 1988; Caon, Capacciolli \\& D'Onofrio 1993;  Young \\& Currie 1994; Graham et al. 1996), and for the bulges of spiral galaxies (Andredakis, Peletier \\& Balcells 1995; Seigar \\& James 1998; Moriondo, Giovanardi \\& Hunt 1998; Khosroshahi, Wadadekar \\& Kembhavi 2000; Prieto et al. 2001; Graham 2001; M\\\"ollenhoff \\& Heidt 2001).  These systems are  not universally described with either an exponential profile or an $r^{1/4}$ law (de Vaucouleurs 1948, 1959), but rather a continuous range of  light profile shapes exist which are well described by the S\\'ersic (1968) $r^{1/n}$ model.  In ellipticals, the shape parameter $n$ from the S\\'ersic model is strongly correlated with the other global properties derived independently of the $r^{1/n}$ model, such as: total luminosity and effective radius (Caon et al.\\ 1993; Young \\& Currie 1994, 1995;  Jerjen \\& Bingelli 1997; Trujillo, Graham \\& Caon 2001), central velocity dispersion (Graham, Trujillo \\& Caon 2001) and also central supermassive black hole mass (Graham et al.\\ 2001).  Additionaly, the spiral Hubble type has been shown to correlate with the bulge index $n$ such that early--type Spiral galaxy bulges have larger values of $n$ than late--type spiral galaxy bulges (Andredakis et al. 1995; Graham 2001).  This correlation arises from the fact that the index $n$ is well correlated with the bulge--to--disc luminosity ratio (B/D; see, e.g. Simien \\& de Vaucouleurs 1986) and this is one of the parameters used to establish morphological type (Sandage 1961).  Given the abundance of observational work and papers now using the S\\'ersic model, it seems timely  that a theoretical study is performed on realistic, analytical models following  the $r^{1/n}$ law.  Structural and dynamical properties of isotropic,  spherical galaxies following $r^{1/n}$ models have already been studied in detail in the insightful paper of Ciotti (1991). However, most  elliptical galaxies and bulges of spiral galaxies are known to be non--spherical objects.  Typically, the mass models which have  been used for the study of triaxial galaxies have followed  analytical expressions which were selected to reproduce the properties of the de Vaucouleurs $r^{1/4}$ profile (e.g. Jaffe 1983; Hernquist 1990; Dehnen 1993), or more recently the modified Hubble law (Chakraborty \\& Thakur 2000). For that reason, previous studies based on these kinds of analytical models, although certainly useful, are however unable to  probe the full range of properties which are  now observed in real galaxies.  Consequently, it  is of importance to know how much  room for  improvement exits in the  study of triaxial objects following the $r^{1/n}$ family of profiles.  Due to the fact that the observed $r^{1/n}$ luminosity profiles cannot be deprojected to yield analytical expression for the spatial density\\footnote{Recently, Mazure \\& Capelato (2002) have provided an exact solution for this, and other related spatial properties, in terms of the Meijer G functions when the  S\\'ersic index $n$ is an integer.}, the $r^{1/n}$ law has been considered less useful for studies of detailed modelling.  An analytical representation (approximation) for the mass density profiles which  accurately reproduces the observed $r^{1/n}$ luminosity profiles would be  of  great interest for simulations of real galaxies. We have therefore derived just such an analytical expression for the mass density profiles  of the S\\'ersic family of models. Our approximation surpasses the accuracy of both the Dehnen models for representing the specific  $r^{1/4}$ profile and also their extension to the double power--law models of Zhao (1997).  In this paper we present a detailed study of how the physical properties of triaxial stellar systems change as a function of the index $n$. An accurate analytical expression for modelling the spatial density is presented  in Section 2.  In Section 3 we explore the   axisymmetric and the non--axisymmetric components of  the potential,  forces and torques associated with a S\\'ersic light distribution. Finally,  by using literature available K--band observations of a sample of 80 spiral galaxies, the physical basis to the $n$--$T$ (or $n$--$B/D$) relation is investigated in Section 4.   \\section[]{The \\lowercase {$r^{1/n}$} model}  The projected, elliptically symmetric S\\'ersic $r^{1/n}$ intensity distribution  $I({\\bf r})$ can  be written in terms of the projected, elliptical radial coordinate $\\xi$ (see details in Trujillo et al. 2001) such that: \\begin{equation} I(\\xi)=I(0)e^{-b_n(\\xi/r_e)^{(1/n)}}, \\end{equation} where $I(0)$ is the central intensity, and $r_e$ is the effective  radius of the projected major--axis. Curves of constant $\\xi$ on the plane of the sky are the isophotes. The quantity $b_n$ is a function of the shape parameter $n$, and is chosen so  that the effective radius encloses half of the total luminosity.  The exact value is derived from  $\\Gamma (2n)$$=$$2\\gamma (2n,b_n)$, where $\\Gamma(a)$ and $\\gamma(a,x)$ are  the gamma function and the incomplete gamma function respectively (Abramowitz \\& Stegun  1964, p. 260). The index $n$ increases monotonically  with the central luminosity concentration of the surface brightness distribution (Trujillo, Graham, \\& Caon 2001).  The total projected luminosity L associated with this model is given by \\begin{equation} L = I(0)r_e^2(1-\\epsilon)\\frac{2\\pi n}{b_n^{\\;2n}}\\Gamma(2n), \\end{equation} where $\\epsilon$ is the ellipticity of the isophotes. For a homologous triaxial ellipsoid, the spatial (deprojected) luminosity density profile $\\nu(\\zeta)$ can be obtained by an Abel integral equation (Stark 1977): \\begin{equation} \\nu(\\zeta)=-\\frac{f^{1/2}}{\\pi}\\int_{\\zeta}^{\\infty} \\left[\\frac{d}{d\\xi}I(\\xi)\\right](\\xi^2-\\zeta^2)^{-1/2}d\\xi , \\end{equation} where $f^{1/2}$ is a constant that depends on the three-dimensional spatial orientation of the object (Varela, Mu\\~noz-Tu\\~noz \\& Simonneau  1996; Simonneau, Varela \\& Mu\\~noz-Tu\\~noz 1998) and $\\zeta$ parametrizes the ellipsoids of constant volume brightness. $f^{1/2}$ equals 1 when the proper axis frame of the object has the same orientation as the observer axis frame (i.e.\\ when the Euler angles between the two frames equal zero).    \\subsection{Mass density  profiles}  Assume a triaxial object whose mass is stratified over ellipsoids with axis ratios a:b:c (a$\\geq$b$>$c) and the x-- (z--) is the long (short) axis (see Fig.\\  1). The symmetry of the problem motivates us to work with ellipsoidal coordinates  where:  \\begin{eqnarray}  x&=&\\zeta  \\sin \\psi \\cos  \\theta \\nonumber\\\\ y&=&\\alpha \\zeta \\sin \\psi \\sin  \\theta \\nonumber\\\\  z&=&\\beta\\zeta \\cos \\psi \\end{eqnarray} and where $\\alpha$=b/a and $\\beta$=c/a. \\begin{figure} \\centerline{\\psfig{figure=MB563fig1.ps,width=6cm}} \\caption{A surface of constant density for the triaxial ellipsoid described in Eq. (5) and (6).} \\end{figure} The mass models considered in this study are the triaxial generalizations of the spherical models discussed in detail by Ciotti (1991). The mathematical singularities present in Eq. 3 were considered and solved by Simonneau \\& Prada (1999, Eq. 16). Substituting Eq. 1 into Eq. 3, letting $\\xi=\\zeta\\cosh\\varphi$, and multiplying by the  mass--to--light ratio $\\Upsilon\\equiv$M/L, we obtain a similar expression to the one found by these authors: \\begin{equation}  \\rho (\\zeta) =  \\frac{f^{1/2}I(0)b_n}{\\pi n r_e^{1/n}} \\Upsilon \\int_0^\\infty e^{-b_n \\left( \\frac{\\zeta\\cosh\\varphi}{r_e}\\right)^{1/n}}(\\zeta\\cosh\\varphi)^{1/n-1}d\\varphi , \\end{equation} with   \\begin{equation} x^2+\\left(\\frac{y}{\\alpha}\\right)^2+\\left(\\frac{z}{\\beta}\\right)^2=\\zeta^2. \\end{equation} The dimensionless mass density profiles $\\widetilde{\\rho}(\\zeta)\\equiv r_e^3/M\\rho (\\zeta)$, where $M$ is the total mass, are shown for different values of $n$ in Fig. 2a. It should be noted that the inner density profile decreases more slowly with increasing radius for systems having lower values of $n$.  The mass density profiles of the $r^{1/n}$ family (Eq. 5) can be extremely well approximated  by the analytical expression: \\begin{equation} \\rho_{app}(\\zeta)=\\frac{f^{1/2}I(0)b_n2^{(n-1)/2n}}{r_en\\pi}\\Upsilon \\frac{h^{p(1/n-1)}K_{\\nu}(b_nh^{1/n})}{1-C(h)}, \\end{equation} where $h=\\zeta/r_e$,  $C(h)=h_1(\\log h)^2+h_2\\log h+h_3$ and $K_{\\nu}$ is the $\\nu$th--order modified Bessel function of the third kind (Abramowitz \\& Stegun 1964, p. 374). In the Appendix A we show the values of the parameters ($\\nu$,p,h$_1$,h$_2$,h$_3$) as function of the index $n$. This approximation contains two exact cases: $n$=0.5 and $n$=1,  and provides relative error less than 0.1\\% for the rest of the cases (Fig 2b) in the radial range 10$^{-3}\\leq\\zeta/r_e\\leq10^{3}$. This approximation surpasses (by a factor of 10$^2$--10$^4$) the expression presented in Lima Neto, Gerbal \\& M\\'arquez (1999).    \\begin{figure} \\centerline{\\psfig{figure=MB563fig2.ps,width=8cm}} \\caption{a) The  dimensionless mass density profiles (see Section 2.1) for the values of $n$=0.5, 1, 2, 4 and 10. b)  the relative error between the  the analytical approximation proposed in Eq. 7 and the exact solution are shown for the previous values of $n$.} \\end{figure}         \\section[]{Non--axisymmetric perturbations due to a triaxial \\lowercase{$r^{1/n}$} structure}   For three different triaxiality mass distributions:  a) spherical ($\\alpha$=$\\beta$=1); b) moderately triaxial ($\\alpha$=0.75, $\\beta$=0.5); c) highly triaxial ($\\alpha$=0.5, $\\beta$=0.25), we have explored, in detail, the non--axisymmetric gravitational field   over the  z=0 plane (i.e. the disc plane when studying spiral galaxies).    \\subsection{Non--spherical component of the gravitational potential in the plane z=0}  We evaluate this quantity by calculating: \\begin{equation} G(r)\\equiv\\frac{\\Phi_2(r)}{\\Phi_0(r)}, \\end{equation} where $\\Phi_2(r)$ and $\\Phi_0(r)$ are the m=2 and m=0 component of the gravitational potential, such that the nth--order term $\\Phi_m(r)$ is evaluated from the gravitational potential on the z=0 plane $\\Phi(r,\\theta)$ by using the Fourier decomposition (see, e.g. Combes \\& Sanders 1981). Gravitational potential and gravitational force expressions are shown on Appendix B.  The profiles of $G(r)$ for different triaxialities and values of $n$ are shown in Fig.\\ 3.  As it is expected, as the triaxiality increases the non--spherical nature of the gravitational field increases. Also, we highlight that for a given triaxiality, smaller values of $n$ (i.e less concentrated mass distribution) give greater  non-spherical gravitational fields.  The maximum non--axisymmetrical behavior of the potential is obtained at radial distances less than 2 $r_e$. This radial distance is also a function of the index $n$, decreasing a $n$ increases and remains quite independent of the triaxiality of the object. For a moderately triaxial object with $n$=1, the non--axisymmetrical component of the potential can vary some 6\\% between r=0 and r=2$r_e$, and varies some 15\\% for our highly  triaxial model.    For an $n$=1 model, and starting from our moderately triaxial case ($\\alpha$=0.75, $\\beta$=0.50),  we increased the value of $\\beta$ to 0.75.  The results are shown in Fig.\\ 3c and reveal that G(r) varied only mildly.  This figure  shows that the non--axisymmetric effect (along the radial distance) in the z=0 plane is mainly due to how the mass of the bulge is distributed  in this plane.    \\subsection{Non--spherical component of radial gravitational forces in the z=0 plane}  The non--spherical component of the radial gravitational forces in the z=0 can be estimated by: \\begin{equation} N(r)\\equiv\\frac{\\partial\\Phi_2(r)/\\partial r}{\\partial\\Phi_0(r)/\\partial r}. \\end{equation} In Fig.\\ 3 the N(r) profiles (Eq. 9) are evaluated for the same cases as was the G(r) profiles. A  remarkable point is that N(r) reaches its maximum value in the radial range 2 $r_e$$<$r$<$4 $r_e$. For a spiral bulge structure, this means that the most important non--axisymmetric effects take place in a zone which is dominated by the disc. As with the G(r) parameter,  stronger distortions occur as the triaxiality increases and the index $n$ decreases. The mechanism which  controls this distortion is basically determined by the mass distribution in the z=0 plane (Fig 3c).  It is noted that the relative (i.e. \\% change) non--axisymmetric effects on the radial forces are larger than the relative distortion on the potential. As an example, for a moderately triaxial structure with $n$=1 the non--axisymmetric component of the radial forces can reach   8\\%.     \\subsection{Torques on z=0 plane}  The torques provoked by the triaxial structures along the angular coordinate are evaluated  around the circle of radius $r_{max}$ where the maximum non--axisymmetrical distortion of the radial forces is produced (i.e. at the peak of the N(r) profile).  Given the gravitational potential $\\Phi(r,\\theta)$ in the z=0 plane, we have at the  radius $r_{max}$: \\begin{equation} \\Pi(\\theta)\\equiv\\frac{F_T(r_{max})}{F_R(r_{max})}, \\end{equation} where $F_T(r_{max})=[\\partial\\Phi(r_{max},\\theta)/\\partial \\theta]/r_{max}$ represents  the  amplitude of the tangential force along the angular coordinate at radius $r_{max}$, and $F_R(r_{max})=(\\partial\\Phi(r_{max},\\theta)/\\partial r)$ is the radial force at this radius. Due to the symmetry of the ellipsoid, the values of $\\Pi(\\theta)$ need only be plotted for one quadrant in the z=0 plane; we use 0$^\\circ$$<\\theta<$90$^\\circ$ (Fig.\\ 3). Depending on the quadrant, $\\Pi(\\theta)$ is either negative or  positive because the sign of the tangential force changes from quadrant to quadrant.  The maximum torque around a circle of radius $r_{max}$  depends on the triaxiality of the object. As the triaxiality increases the maximum torque tends to be closer to the major axis -- as one would expect. The position of this peak is quite independent of the value of  $n$.  The absolute value of the torque for any given triaxiality increases  as $n$ decreases.   For our highly triaxial bulge, $\\Pi(\\theta)$ ranges from 0.17 ($n$=10) to 0.24 ($n$=1),  which would be considered a \"bar strength\" class of 2 in the classification  scheme of Buta \\& Block (2001). In the case of our moderately triaxial object, the maximum  absolute value of $\\Pi(\\theta)$ ranges between 0.06  and 0.09. These values correspond to a ``bar strength'' class of 1.  Thus, even a  moderately triaxial bulge is capable of provoking non--negligible  torques on a disk -- that is to say, a bar is not neccessarily required.   A detailed study separating the torque contribution from both bars and  bulges would of course be of interest, and it is our intention to add a range of bar potentials to our models in the future.  As with the previous parameters, for the range of triaxialities investigated and a given $n$, varying the mass distribution along the  z--axis (i.e. varying the triaxiality parameter $\\beta$) only results in a  slight change to $\\Pi(\\theta)$  (see Fig. 3c). For a spherical distribution all above parameters are 0.  \\begin{figure} \\centerline{\\psfig{figure=MB563fig3.ps,width=8cm}} \\caption{The parameters G(r), N(r) and $\\Pi(\\theta)$ are shown for different values of $n$ and triaxiality: a) First Row: G(r), N(r) and $\\Pi(\\theta)$ for a moderately triaxial object and different $n$; b) Second Row: G(r), N(r) and $\\Pi(\\theta)$ for a highly triaxial structure; c) Third Row: G(r), N(r) and $\\Pi(\\theta)$  for three moderately triaxial objects with $n$=1 and the same axis ratio along the y and x axis.}  \\end{figure} ",
        "conclusions": "The main results of this work are the following:  a) We have generalised the analysis of the physical properties of spherical stellar systems following the $r^{1/n}$ luminosity law   to a homologous triaxial distribution. The density distribution, potential, forces and torques are evaluated and compared with the spherical case when applicable (Ciotti 1991). An extremely accurate analytical approximation (relative error less than 0.1\\%) for  the mass density profile is provided.  b) We derive an exact expression showing how the central potential decreases as triaxiality increases. We also show that for a fixed triaxiality, as the index $n$ decreases the  non--axisymmetrical effects in the z=0 plane increase. Even for a moderately triaxial object, the non--axisymmetrical component of the potential and the radial forces are not negligible for small values of $n$. These components can range from  6 to 8\\%, respectively, compared to the value of the spherical component. For our highly triaxial model, they can range over some 20\\%.   c) The non--axisymmetrical effects in the disc plane due to the bulge structure are strongly reduced when an axisymmetrical disc mass is added.   For this reason, bulges with smaller values of $n$ appear unlikely to produce any significant non--axisymmetrical effect on their disc, which is typically 10 to 100 more times more massive than the bulge. In this regard, the $B/D$ mass ratio and the triaxiality of the bulge are more important, that is, can dominate over the effects of small $n$.  d) The correlation found between $n$ and the $B/D$ luminosity ratio found in spiral galaxies is  explained here not as a consequence of the interplay between the bulge and the disc,  but due to the strong correlation between $n$ and  $M_{T}(bulge)$, and between $M_{T}(bulge)$ and  $B/D$. Also, K--band data do not support the idea that the $B/D$ luminosity ratio can be preferred over $M_{T}(bulge)$ as an indicator to establish galaxy morphological type (T). Both parameters present equally good correlations with galaxy type T."
    },
    "0201/astro-ph0201216_arXiv.txt": {
        "abstract": "We estimate the number of ${\\mathrm z} \\simeq 7$ quasars that will be discovered in the Large Area Survey (LAS) element of the UKIRT Infrared Deep Sky Survey (UKIDSS). The LAS will cover 4000 sq degs of the northern sky to $K=18.4$, which is 3 mag. deeper than 2MASS. The Sloan Digital Sky Survey has extended the quasar redshift limit to ${\\mathrm z}=6.3$. We demonstrate that to reach higher redshifts ${\\mathrm z}\\sim7$, when Ly$\\alpha$ has passed through the $z$ band, combinations of standard broad-band filters such as $zJH$ and $zJK$ are ineffective. Instead the wavelength range between $z$ and $J$ must be exploited. We introduce the $Y$ passband $0.97-1.07\\mu{\\mathrm m}$ for this purpose. High-redshift quasars up to redshift ${\\mathrm z}=7.2$ can be selected from a $iYJ$ or $zYJ$ two-colour diagram, as bluer than L and T dwarfs. ",
        "introduction": "Over the past two years the Sloan Digital Sky Survey (SDSS) has had considerable success in searches for quasars of very high redshift ${\\mathrm z}>5$ (e.g. Schneider, these proceedings). Currently the highest redshift recorded is ${\\mathrm z}=6.3$ (Fan et al., 2001a). Perhaps the most important motivation for searching for sources at even higher redshifts is the possibility of detecting and studying the epoch of reionisation. Analysis of the optical depth blueward of the Ly$\\alpha$ emission line in the spectrum of the ${\\mathrm z}=6.3$ quasar indicates that we may be on the threshold of this event (Becker et al., 2001). The detection of quasars of higher redshift could probe the neutral IGM during or even before the epoch of reionisation. Although the optical depth in Ly$\\alpha$ is too high to provide a useful measure of the neutral fraction in the IGM (e.g. Madau, these proceedings), other species, observable with NGST, might provide a measure. Also other probes of the IGM have been suggested (e.g. Loeb and Rybicki, 2000).  \\begin{figure} \\plotone{sjwdur1.eps} \\caption{Two-colour $zJK$ diagram showing colours of stars O to M (filled circles), and L and T brown dwarfs, open and filled squares respectively. All magnitudes are on the Vega system. The chains show model quasar colours $5<{\\mathrm z}<8$, $\\Delta {\\mathrm z}=0.1$, with three different continuum slopes. This two-colour diagram is ineffective for selecting quasars in the redshift range $6<{\\mathrm z}<7$, and other combinations of $zJHK$ are no better. Quasars ${\\mathrm z}>7$ could be selected by identifying objects $z-J>3.2$, i.e. in the box shown. But to reach $J=19$ requires a depth $z=22.2$, whereas the SDSS only reaches $z=19.9$.} \\end{figure}  \\begin{figure} \\plotone{sjwdur2.eps} \\caption{Two-colour $iYJ$ diagram showing colours of stars O to M (filled circles), and L and T brown dwarfs, open and filled squares respectively. All magnitudes are on the Vega system. The chains show model quasar colours $5<{\\mathrm z}<8$, $\\Delta {\\mathrm z}=0.1$, with three different continuum slopes. Any objects $i-Y>3$, in the selection box shown, are candidate quasars in the redshift range $5.8<{\\mathrm z}<7.2$. Because the SDSS reaches $i=22.0$ this works to a limit $Y=22.0-3.0=19.0$. In fact the UKIDSS LAS will reach $Y=20.5$ so deeper optical data will allow the discovery of fainter high-redshift quasars.}  \\end{figure} ",
        "conclusions": ""
    },
    "0201/astro-ph0201020_arXiv.txt": {
        "abstract": "The  holographic entropy bound is used to estimate the  quantum-gravitational discreteness of inflationary perturbations. In the context of    scalar inflaton perturbations produced during standard slow-roll inflation,  but assuming that horizon-scale perturbations ``freeze out'' in discrete steps   separated by one bit of total observable entropy, it is shown that  the Hilbert space of a typical horizon-scale inflaton perturbation is equivalent to that of  about $10^5$ binary spins--- approximately  the inverse of the final scalar metric perturbation amplitude, independent of other parameters. Holography thus suggests that in a broad  class of fundamental theories, inflationary perturbations carry  a  limited amount of information (about $10^5$ bits per mode) and should therefore  display discreteness not predicted by the standard field theory. Some manifestations  of this discreteness may be observable  in cosmic background anisotropy. ",
        "introduction": "The origin of cosmological perturbations now appears to be well understood  from  the quantum theory of fields in curved spacetime \\cite{Starobinsky:1979ty,Hawking:1982cz,Guth:1982ec,Bardeen:1983qw,Starobinsky:1982ee,Halliwell:1985eu,Grishchuk:1993ds}. They originate during inflation as zero point fluctuations of the quantum modes of various fields--- the inflaton   giving rise to scalar perturbations, the graviton to tensor perturbations. The field quanta in the original fluctuations convert to classical perturbations as they pass through the de Sitter-like event horizon;   they are then parametrically  amplified by an  exponential factor during the many subsequent e-foldings of inflation, creating an enormous number of coherent quanta in phase with the original quantum seed perturbation. From the classical point of view,  the quantum fluctuations create perturbations in the classical gravitational gauge-invariant potential  $\\phi_m$\\cite{bardeen1980}, leading to  observable background anisotropy and  large scale structure\\cite{cobeDMR,bennett,gorski,boom,pryke,halverson,max,bond01}. All stages of this process are under good calculational control, even the conversion of quantum to classical regimes \\cite{Grishchuk:1989ss,Grishchuk:1990bj,Albrecht:1994kf,Lesgourgues:1997jc,Polarski:1996jg,kiefer}. The phase and amplitude of the large-scale classical perturbation modes observed today are a direct result of the quantum field activity during inflation; indeed, the pattern of microwave anisotropy on the largest scales corresponds to a faithfully amplified image  of microscopic field configurations as they froze out during inflation. Roughly speaking, each hot or cold patch on the  sky derives originally from about one quantum.  The standard  calculation  of these processes\\cite{Mukhanov:1992me,lythriotto} uses a semiclassical  approximation: spacetime is assumed to be classical (not quantized), and the perturbed fields (the inflaton and graviton) are described using relativistic quantum field theory, essentially (in the limit of  free massless fields) an infinite collection of quantized harmonic oscillators. The Hilbert space of this system is infinite, so although the fields are quantized, they are continuously variable functions that can assume any values. The creation of the particles can be viewed as an effect of the nonadiabatic expansion, and the ``collapse of the wavefunction'' (in this case, ``freezing out'') can be described as a unitary quantum process of state squeezing. The theory generically predicts random-phase gaussian noise with a continuous spectrum determined by the parameters of the inflaton potential.  In this approximation, there is no telltale  signature of    quantum discreteness in the  observable classical remnant--- the anisotropy of the background radiation. Although sky maps contain images of ``single quanta,'' their spectrum is continuous and the amount of information  is in principle infinite.   It has always been acknowledged that this description is incomplete, and will be modified by including a proper account of spacetime quantization. Although the fundamental theory of such quantization is not known, a ``holographic entropy bound'' already  constrains with remarkable precision the total number of fundamental quantum degrees of freedom. The complete Hilbert space of a bounded volume is finite and discrete  rather than infinite and continuous, limiting the range of accessible configurations in any region to a definite, calculable number.   In particular this limit applies to   inflaton quanta   collapsing into  classical metric perturbations.   This paper  uses the holographic entropy bound to  estimate   where   field theory breaks down in the inflationary analysis, the effective dimension of the Hilbert space for the observable perturbations when they freeze out, the maximum amount of information contained in the perturbations, and the level at which quantum-gravitational discreteness  appears in cosmic background anisotropy. The main result is that in fundamental theories where the holographic entropy bound arises from discrete fundamental eigenstates, the amount of information in the anisotropy is remarkably limited: it can be  described with only about $10^5$ bits per sky-harmonic mode, implying that the  perturbations should be pixelated in some way. In principle, this effect may be observable, and provide concrete data on the discrete elements or eigenstates of quantum gravity. ",
        "conclusions": ""
    },
    "0201/hep-ph0201273_arXiv.txt": {
        "abstract": "Neutrinos from far away sources annihilating at the Z resonance on relic neutrinos  may give origin to the ultrahigh-energy cosmic rays. Here we present predictions of this mechanism with relic neutrinos lighter than 1 eV, which do not gravitationally cluster. We show that  not only the super GZK events, but the ``ankle\" and all  events above it can be accounted for. Most primaries above the ankle are predicted to be  nucleons up to $10^{20.0}$~eV and  photons at higher energies. We also find an accumulation at the GZK cutoff energy, a hint of  which can  be seen in the data. ",
        "introduction": "\\label{sec-1}  The existence of ultrahigh-energy cosmic rays (UHECR) with energies above the Greisen-Zatsepin-Kuzmin (GZK) cutoff~\\cite{gzk}  of about $5\\times 10^{19}$ eV, presents an outstanding problem~\\cite{data}.  Nucleons and photons with those energies have short attenuation lengths and could only come from distances of 100 Mpc or less~\\cite{50Mpc,40Mpc}, while plausible astrophysical sources for those energetic particles are much farther away.  An elegant and economical solution  to this problem, proposed by T. Weiler~\\cite{weiler}, consists of the production of the necessary photons and nucleons close to Earth, in the annihilation at the $Z$-resonance of  ultrahigh-energy neutrinos coming from remote sources, $\\nu_{\\rm UHE}$, and relic background neutrinos. These events were named ``$Z$-bursts\"  by T. Weiler. One of the most appealing features of the ``$Z$-bursts\" scenarios is that the energy scale of $10^{20-21}$~eV, at which the unexpected events have been detected, is generated naturally given the possible mass range of relic neutrinos. The $Z$-resonance occurs when the energy of the incoming $\\nu_{\\rm UHE}$ is $E_{\\nu_{\\rm UHE}}= E_{Res}$, \\begin{equation} E_{Res} = \\frac{M_Z^2}{2~m_\\nu} , \\label{Eres} \\end{equation} where $m_\\nu$ is the mass of the relic neutrinos. This is the new cutoff of the UHECR energy in these models. It depends inversely on the mass of the relic neutrinos.  Since arguments of structure formation in the Universe show  $m_\\nu <$ few eV, then $E_{\\nu_{\\rm UHE}} > 10^{21}$ eV, precisely above the GZK cutoff, as needed.  In this paper we concentrate on a particular ``Z-bursts'' scenario~\\cite{gk1,gk2}, in which the relic neutrinos are lighter than 1 eV. These lighter neutrinos, contrary to those in the original scenario, cannot be gravitationally bound, they have a constant density over the whole Universe. In particular, we concentrate on relic neutrinos with mass compatible with Super-Kamiokande results, if neutrino masses are hierarchical (however the results we obtain apply to heavier relic neutrinos, while light enough to not cluster). Super-Kamiokande has provided a strong evidence for the oscillation in atmospheric showers of  two neutrino species with masses $m_1$, $m_2$ and  $\\delta m^2 $ = $m_1^2-m_2^2$ = $(1-8) \\times 10^{-3}$ eV, consisting mostly of about equal amounts of $\\nu_\\mu$ and another flavor eigenstate neutrino, $\\nu_\\tau$ or a sterile neutrino~\\cite{SK}.  If neutrino masses are hierarchical, as those of the other leptons and quarks, then the heavier of the two oscillating neutrinos, call it $\\nu_{\\rm SK}$, has a mass $m_{\\rm SK} =\\sqrt{\\delta m^2} \\simeq 0.07$ eV. With $m_\\nu = m_{\\rm SK}$, the new UHECR cutoff becomes \\begin{equation} E_{Res} \\simeq 0.6 \\times 10^{23}~eV . \\label{Eres2} \\end{equation} Due to the large multiplicity of the Z-decays, after energy losses in the propagation (as shown in detail here) this value of $E_{Res}$ predicts many super GZK UHECR events (many more than with eV relic neutrino masses).  We agree with Farrar and Piran~\\cite{farrarpiran}, who have argued that any mechanism accounting for the events beyond the GZK cutoff should also account for the events down to the ankle, including their isotropy and spectral smoothness. We show here that the  model we consider can account for the ankle and all the events above it if the  position of the ankle is close to that measured by AGASA, $E_{ankle} = 10^{19.0}$~eV (see \\cite{naganowatson}, in particular Table V, and references therein). Moreover, a  reliable prediction of the model is that most primaries above the ankle should be nucleons  up to about $10^{20.0}$~eV and photons at higher energies. We also find that  nucleons do accumulate at the GZK cutoff energy, which could account for the hint of a slight accumulation seen in the data. Photons become dominant at energies higher than  $10^{20}$~eV in our model. So these photons are all above the threshold energy (which is about $ 5 \\times 10^{19}$~eV) to pair produce in the Earth's magnetic field (which should generate a certain amount of north-south asymmetry in the arrival direction distribution).  Let us return to the issue of the isotropy of the events above the ankle, i.e., the absence of a correlation with the galactic halo. Because the relic neutrinos we assume do not gravitationally cluster, the isotropy of the events reflects the isotropy of the ultrahigh-energy neutrino sources. In particular, relic neutrinos of mass $m_{\\rm SK}$ require a large flux of neutrinos with energies of the order of $10^{23}$~eV. It is unlikely that active galactic nuclei~\\cite{wb}, neutron stars~\\cite{olinto1}, or other astrophysical sources could produce such a high energy flux of ultrahigh-energy neutrinos. Topological defects~\\cite{berez}, or superheavy relic particles~\\cite{gk2}, could instead easily generate the requisite flux of primary neutrinos (but there are still problems with these sources). For example, with unstable superheavy relic particles, which form part of the cold dark matter, decaying mostly into neutrinos~\\cite{gk2},  the directions of UHECR could map the distribution of parent particles (which should coincide with the distribution of cold dark matter) at large redshifts.  This is because the initial energy of the $\\nu_{\\rm UHE}$ produced in the decay needs to be redshifted to the energy of the ``$Z$-burst\" in its way to the Earth. Directional clustering, evident in the existing data~\\cite{clustering}, would then reflect the distribution of matter at a certain red shift determined by this process of ``cosmological filtering''. Thus, absence of directional correlations with the galactic halo, as well as  directional clustering, could be easily accommodated~\\cite{gk2}.  Besides, $\\nu_{\\rm UHE}$ produced by  unstable superheavy relic particles would have a spectrum opposite to an astrophysical spectrum, growing rapidly with energy, up to a sharp cutoff at an energy of the order of the parent particle mass. This spectrum has almost no neutrinos at low energy where bounds exist~\\cite{gk2,sigl}. Most bounds on  ``$Z$-bursts\" models (see for example \\cite{wb}) assume that the $\\nu_{\\rm UHE}$ have a typical astrophysical spectrum, decreasing with energy as $E^{-\\gamma}$, with $\\gamma$ a number of order one.  These bounds do not hold if the $\\nu_{\\rm UHE}$ spectrum has a very different energy dependence. However, a model for these parent particles is arguably difficult to obtain~\\cite{CDF,Uehara}. Moreover, the EGRET bound on the diffuse low-energy gamma ray flux resulting from the ``$Z$-bursts\" imposes important constraints \\cite{EGRET}, which might rule out heavy particles decaying mostly into neutrinos as sources \\cite{Berezinsky}.  In the next section we present our simulations and the resulting spectrum of UHECR. We would like to point out that the main result of this paper, the fact that ``$Z$-bursts\" can account for the ankle and all events above it, does also hold for larger relic neutrino masses, for which the problem of the sources becomes less severe. In fact even if we used $m_{\\rm SK}$  here, our considerations apply with trivial changes to other masses for which relic neutrinos are too light to gravitationally cluster. As the relic neutrino mass increases, all the features in the spectrum we find here should move progressively to lower energies. ",
        "conclusions": "In this paper we considered  a particular ``Z-bursts'' scenario~\\cite{gk1,gk2}, in which the relic neutrinos are lighter than 1 eV, and thus do not gravitationally cluster. Using in particular 0.07 eV relic neutrinos, we have shown that ``$Z$-bursts\"  may account  not only for the super GZK events, but for the ``ankle\" and all UHECR events above it, including their isotropy and spectral smoothness. In our simulation we found the ``ankle\"  close to $E_{ankle} = 10^{19.0}$~eV as measured by AGASA. A reliable prediction of the model is that most primaries above the ankle should be nucleons up to about $10^{20.0}$~eV and photons at higher energies. Moreover, the nucleons do accumulate at the GZK cutoff energy, which could  account for the hint of a slight accumulation seen in the data. The model predicts a new cutoff, which with 0.07 eV relic neutrinos is at $E_{Res} \\simeq 0.6 \\times 10^{23}~eV$.  We have not included the effect of extragalactic magnetic fields, thus for the predictions of this paper to be true these fields should be sufficiently small, probably  smaller than $10^{-9}~G$.  Finally let us comment on recent related papers. Photon and nucleon spectra very similar to those presented here are shown in Fig. 2a of Ref. \\cite{kalashev} corresponding to ``$Z$-bursts\" with 0.1 eV relic neutrinos and a different model for the  distribution of ``$Z$-bursts\" with redshift up to $z_{max}= 3$. This model satisfies the EGRET bound on low energy photons (even if with astrophysical sources emitting only UHE neutrinos).  This reassures us that the result we present here is robust. The EGRET bound has been computed for various ``$Z$-bursts\" scenarios~\\cite{kalashev,Yoshida}, including the particular one we concentrated on here~\\cite{EGRET}, which seems to work well with sources emitting only UHE neutrinos. The most serious problem with these sources may be the electroweak cascading of the UHE neutrinos produced in the decays, as recently claimed in~\\cite{Berezinsky}.  Events above the ankle have been previously fitted with ``$Z$-bursts\" products plus an additional hypothetical component of galactic or extragalactic protons in Ref.~\\cite{Fodor}, varying the relic neutrino mass, in an attempt to provide a determination  of this mass using UHECR data. However the normalization and slope of the extra proton flux and the normalization of the nucleon flux from ``$Z$-bursts\" (photons from ``$Z$-bursts\" were neglected) were taken as fitting parameters. This procedure does not make clear if it is actually the ``$Z$-bursts\" which account for the change of slope at the ``ankle\" and for the events above the ``ankle\". Moreover, with this procedure~\\cite{Fodor}, there is no prediction  of the position of the ankle, since this is one of the parameters to be fitted by the sum of the mentioned two fluxes. Here we took instead the measured flux below the ankle, with the measured slope and normalization. We considered it to be of galactic origin, as the correlation with the galactic center of the arrival directions measured by AGASA around $10^{18}~eV$~\\cite{correlation} seems to indicate. Then we added to this fixed flux the new component due to ``$Z$-bursts\", which led us to find a prediction for the position of the ankle, and the normalization of the flux of primaries due to ``$Z$-bursts\".  We believe our flavor of ``$Z$-bursts\" provides  a plausible explanation to the puzzle of ultrahigh-energy cosmic rays, not only for the super GZK events, but for the ``ankle\" and all  events above it.  \\vspace{5 mm}    This work was supported in part by the US Department of Energy grant DE-FG03-91ER40662, Task C. G.V. was also supported by an award from Research Corporation. We thank A. Kusenko and S. Nussinov for many valuable discussions, and P. Biermann and E. Roulet for comments and suggestions."
    },
    "0201/astro-ph0201093_arXiv.txt": {
        "abstract": "We report on mid-infrared observations of the compact stellar cluster located in the proximity of SGR 1900+14, and the radio/X-ray position of this soft-gamma repeater. Observations were performed in May and June of 2001 when the bursting source was in an active state. At the known radio and X-ray position of the SGR we did not detect transient mid-IR activity, although the observations were performed only hours before and after an outburst in the high-energy band. ",
        "introduction": "Recent deep, high-resolution multi-wavelength timing observations of SGRs led to significant progress in our understanding of these enigmatic sources [1, 2]. Key goals in current SGR studies include the identification of their counterparts at long-wavelengths and a better understanding of their past and future evolutionary states (e.g., [3]). Of particular interest is their possible relation to the class of anomalous X-ray pulsars [4].  From the point of view of ground-based astronomy, among the known four (perhaps five) soft gamma-ray repeaters [1] SGR 1900+14 has the advantage that it (or better, its error box) is observable from the northern as well as the southern hemisphere. This increases the opportunities to monitor this source in the optical/infrared bands whenever it is in an active state. During its recent activity cycle in spring/summer 2001 we observed the SGR 1900+14 error box with the ESO 3.6-m telescope using the newly commissioned TIMMI~2 mid-infrared camera. The campaign covered the position shortly after and before an outburst in the high-energy band. ",
        "conclusions": "The issue of whether or not SGR 1900+14 and the compact stellar cluster its proximity are physically related is crucially linked to the very uncertain distances. Are these objects located at the same distance or not? In this particular case, this question could be answered if a better understanding of the measured extinction toward the M supergiants and toward the SGR is achieved. The former might profit from deeper mid-infrared observations, the latter might gain from the recent discovery [27] of a persistent dust-scattered X-ray halo around the quiescent X-ray counterpart of the SGR, which could lead to a direct measurement of the extinction by the scattering dust along the line of sight [28].  The non-detection of any signal from the SGR might not be surprising if there is no accretion disk at all around the burster, as indicated by the observations performed to date [20]. But should one expect to detect any non-gamma-ray signal from the burster within hours of a high-energy outburst? This question remains to be addressed by further theoretical studies, but we note that there are now two cases where mid-IR observations were performed only a few hours after a SGR outburst (the other case is SGR 1820--20 [21]) and no signal from the burster or the ambient interstellar medium was detected down to a flux density limit of a few mJy. Future observations need to further push the sensitivity limit, and also sample more closely in time. Truly simultaneous coverage would require robotic observations, similar to those used in the search for prompt optical emission [29, 30].   \\begin{theacknowledgments} The authors thank Jochen Greiner and U.R.M.E. Geppert (both AIP Potsdam, Germany) for valuable comments on the manuscript. \\end{theacknowledgments}"
    },
    "0201/astro-ph0201436_arXiv.txt": {
        "abstract": "The EPIC detectors on \\xmm provide the most sensitive broad band ($\\sim 0.3-12$~keV) X-ray spectra to date. Observations of 6 Seyfert 1 galaxies, covering a wide luminosity range, are examined with the aim of identifying the primary X-ray continuum and constraining superimposed emission and absorption features. A soft excess $\\it emission$ component is seen in every case, but with a spectral form which differs markedly with luminosity across our sample. Current interpretations of the soft excess range from intrinsic thermal emission from the accretion disc to reprocessing of harder radiation absorbed in the disc skin. Visual examination of the broad-band EPIC spectra suggest that the luminosity trend in the observed spectral profiles may be governed primarily by differences in the line-of-sight absorption. In that case the underlying continuum could have a common form across the sample. Examination of spectral features in the Fe K band confirm the common presence of a narrow emission line at $\\sim$ 6.4~keV. Modeling of the EPIC spectra above $\\sim$ 7~keV is shown to be critical to quantifying (or confirming) a broad Fe K line in at least some cases. ",
        "introduction": "For more than a decade after powerful X-ray emission was established as a common property of AGN (\\cite{kpounds-C2:elvis78}, Pounds 1977) little insight on the nature of the X-ray source was gained from spectra which were well-modelled by a featureless power law continuum of photon index $\\Gamma$~$\\sim1.5-2$. The increased spectral bandwidth of the combined ME and LE detectors on \\exosat yielded evidence that the spectrum of many AGN steepened below $\\sim0.2-12$~keV, the so-called 'soft excess' (Arnaud 1985, Turner \\& Pounds 1989). Since that time new X-ray satellites have brought major improvements in spectral resolution, sensitivity and bandwidth, providing X-ray spectra of increasing complexity (and diagnostic potential). The particular strengths of each new mission has led to new areas of research, with the high energy response of \\ginga establishing the widespread importance of `reflection' (Nandra \\& Pounds 1994), \\rosat showing absorption in \\los photoionised matter to be surprisingly strong (Turner \\et 1993), \\asca resolving discrete spectral features, including a broad fluorescent iron K line (Tanaka \\et 1995, Nandra \\et 1997) and \\sax providing uniquely broad bandwidth data.  \\chandra and \\xmm have now taking this evolution a large step further, with high-throughput grating spectra providing a qualitative improvement in spectroscopic detail. The most striking results to date have been in the detection of complex absorption and emission spectra in the soft X-ray band (e.g. Kaspi \\et 2001) where the 1- or 2-stage photoionisation model of the `warm absorber' (e.g. Reynolds 1997) has been shown to be an inadequate description. Instead, the \\xmm and \\chandra grating spectra have shown the (often outflowing) material illuminated by the central X-ray source to exhibit a remarkably wide range of ionisation stages (Behar \\et 2001). Interpretation and modelling of these data is already promising a dramatic improvement in mapping the structure, content and dynamics of matter in the nuclear region of Seyfert galaxies. Conversely, modelling the temperature and ionisation of this gas will provide a unique measure of the total ionising flux, including the often dominant XUV component invisible from Earth.  To date, less progress has been made in utilising the diagnostic potential of the broader-band spectra from the EPIC cameras on \\xmm; to briefly explore that potential is the purpose of the present paper which brings together EPIC spectra from 6 Seyfert galaxies observed early in the mission and covering the (2-10 keV) X-ray luminosity range $10^{43}-10^{45}$erg sec$^{-1}$. ",
        "conclusions": "The most striking aspect of the range of EPIC spectra displayed in Figure 1 is the presence in every case of a strong soft emission component which causes the X-ray continuum to turn up sharply below $\\sim$~0.7 keV in the lower luminosity sources (SSX), and is evident as a gradual up-turn (GSX), below $\\sim$~3 keV for the high luminosity sources. The form and location of the SSX is strongly suggestive of a link with K-shell oxygen, though - unsurprisingly - the lower resolution EPIC data can not distinguish between the alternative descriptions (relativistically broadened line emission, Sako \\et 2001; complex absorption, Lee \\et 2001) currently being advocated for MCG-6-30-15 and Mkn 766. RGS spectra do confirm the absence of strong absorption in the higher luminosity sources in our sample, Mkn 509 and PKS 0558-504. This leads us to speculate that a broad soft emission component extending up to $\\sim$~3~keV (arising by Comptonisation of cool disc photons) may be common to our sample, but is obscured by increasingly strong absorption from ionised gas for the lower luminosity Seyferts NGC 5548, Mkn 766 and MCG-6-30-15. We emphasise the complementary strengths of broad-band data from EPIC and higher resolution spectra from the RGS in providing the input for improved modelling of the primary X-ray continuum and - consequently - of absorption and secondary emission features that carry unique information on the X-ray emission mechanism and nuclear environment in Seyfert galaxies.  One such secondary emission feature, the narrow Fe K emission line at 6.4 keV, is found to be a common property across our sample (apart from the highest luminosity object PKS 0558-504). The EW of $\\sim$~50--100 eV is consistent with reflection from cold, distant matter subtending a solid angle of 1--2 $\\pi$ steradians at the hard X-ray source. Line profile and variability studies offer the promise of using this emission line to probe dense cold matter from the outer disc to the molecular torus.  We suggest the broad Fe K line may not be as common, or as strong, as it appeared to be from earlier \\asca observations. A key factor is the need to correctly model the spectrum above $\\sim$~7 keV, where \\xmm has improved - though still limited - sensitivity. Also, as spectral fits to the MCG-6-30-15 data show, precise calibration of X-ray optics and detectors is critical to validating measurements of the important broad Fe K line."
    },
    "0201/astro-ph0201371_arXiv.txt": {
        "abstract": "In this paper, we present moderately-high resolution ($\\sim 65\\;\\kmsm$) spectroscopy, acquired with ESI on Keck II, of 11 ultraluminous infrared galaxies at $z < 0.3$ from the IRAS 1 Jy sample.  The targets were chosen as good candidates to host galaxy-scale outflows, and most have infrared luminosities dominated by star formation.  We use a $\\chi^2$ minimization to fit one- to three-component profiles to the \\ion{Na}{1} D interstellar absorption doublet in each object.  Assuming that gas blueshifted by more than 70 \\kms\\ relative to the systemic velocity of the host is outflowing, we detect outflows in $73\\%$ of these objects.  We adopt a simple model of a mass-conserving free wind to infer mass outflow rates in the range $\\dot{M}_{\\mathrm{tot}}$(H)$\\;= 13-133\\;\\smpym$ for galaxies hosting a wind.  These values of $\\dot{M}_{\\mathrm{tot}}$, normalized to the corresponding global star formation rates inferred from infrared luminosities, are in the range $\\eta \\equiv \\dot{M}_{\\mathrm{tot}} / \\mathrm{SFR} = 0.1-0.7$.  This is on average a factor of only 10 less than $\\eta$ from recent measurements of nearby dwarfs, edge-on spirals, and  lower-luminosity infrared galaxies.  Within our sample, we conclude that $\\eta$ has no dependence on the mass of the host (parameterized by host galaxy kinematics and absolute $R$- and $K^{\\prime}$-band magnitudes).  We also attempt to estimate the average escape fraction $\\langle f_\\mathrm{esc} \\rangle \\equiv \\sum \\dot{M}_\\mathrm{esc}^i / \\sum \\dot{M}_{\\mathrm{tot}}^i$ and ``ejection efficiency'' $\\langle\\delta\\rangle \\equiv \\sum \\dot{M}_\\mathrm{esc}^i / \\sum \\mathrm{SFR}^i$ for our sample, which we find to be $\\sim 0.4-0.5$ and $\\sim 0.1$, respectively.  The complex absorption-line properties of Mrk 231, an ultraluminous infrared galaxy which is optically classified as a Seyfert 1, are discussed separately in an appendix. ",
        "introduction": "\\label{intro}  Large-scale galactic outflows, energized by stellar winds and supernovae ejecta or a central AGN, are ubiquitous in the local universe and at high redshift \\citep*{pet01,fbb02}.  Theoretical work, both analytical and numerical, suggests that these outflows may be important in galaxy formation \\citep*[e.g.][]{ds86,sfb00,sb01,std01}.  They also play a role in galactic evolutionary processes, especially the regulation of star formation by mechanical feedback \\citep*[e.g.][]{ocm01} and the enrichment of galactic halos.  The expulsion of metals from galaxies by galactic winds may be able to reproduce the color-magnitude (or mass-metallicity) relation of ellipticals, since more massive galaxies, with larger gravitational potentials, could retain mass and metals more easily in this scenario \\citep{larson74,vader86,franx90,kc98}.  Besides their impact on individual galaxies, galactic winds may contribute significantly to the chemical and thermal evolution of the universe by enriching and heating the intergalactic medium \\citep*[e.g.][]{nath97,aguirre01,mfr01}, which is known to contain substantial amounts of metals even at high redshift and low density \\citep[e.g.][]{cs98}.  Winds are also a possible mechanism for the preheating and enrichment of the intracluster medium in groups and clusters of galaxies \\citep*[e.g.][]{dfj91,pcn99}.  Finally, these outflows could be responsible for the damped Ly$\\alpha$ or strong \\ion{Mg}{2} absorption systems observed in quasar spectra \\citep[e.g.][]{efsg00,chen01,bond01}, and may contribute to reionization by opening a path for Lyman continuum photons to escape star-forming regions \\citep*{mhr99,dsf99,spa01,pet01}.  The observational dataset on galactic-scale outflows is slim, however, apart from studies of the local universe ($z<0.1$) and high-redshift galaxies [Lyman-break galaxies at $z\\sim3$ \\citep{pet00,pet01} and a few gravitationally-lensed objects \\citep{fbb02}].  In this paper, we present and discuss the results of a moderately-high resolution ($R\\sim4600$, or $\\Delta v \\sim 65\\;\\kmsm$) spectroscopic study of winds in ultraluminous infrared galaxies (ULIGs) with redshifts of $0.04<z<0.27$.  ULIGs, which by definition have $L_{\\mathrm{IR}}>10^{12}L_{\\sun}$\\footnote{$L_{\\mathrm{IR}}=L(8-1000)\\;\\micron$, computed using the single-temperature dust-emissivity fit to all four IRAS flux bands given in \\citet{per87} \\citep[see also][]{sm96}, $H_0 = 75\\;\\kmsm\\;\\mathrm{Mpc^{-1}}$, and $q_0 = 0$.}, contain massive starbursts and/or AGN \\citep*{vks99b,lvg99}, and have been associated with large-scale outflows at low redshift \\citep*[e.g.][]{ham90,kvs98}.  ULIGs, especially those whose energy output is dominated by stars, are important sites of obscured star formation at low $z$ \\citep{sm96}.  There are suggestions from infrared and submillimeter counts that there is a strong increase in the number density of ULIGs with $z$ (IRAS: \\citealt{ks98}; ISO: \\citealt{kawara98,puget99,matsuhara00,efsa00,serjeant01}; SCUBA: \\citealt*{sib97}; \\citealt{hughes98,blain99,eales99}; \\citealt*{bcs99}), implying that ULIGs contain a substantial fraction of high-$z$ star formation (although this is not necessarily true; see \\citealt*{tbg99}).  ULIGs may also be a highly reddened and luminous subset of the UV-selected Lyman-break galaxies \\citep[e.g.][]{sand99, tbg99}, which typically produce outflows \\citep{pet01}.  Absorption-line spectroscopy has proven to be an effective method for probing various phases of the ISM at low and high redshift: the warm neutral ($T\\la10^4$ K) component, using \\ion{Na}{1} D [\\citealt{phillips93} (NGC 1808); \\citealt{heck00}]; the warm ionized ($T\\sim10^{4-5}$ K) component, using UV lines [\\citealt{lequeux95} (Mrk 33); \\citealt{hl97} (NGC 1705); \\citealt{sahu97} (NGC 1705); \\citealt{kunth98,gd98,pet01}]; and the hot coronal ($T\\sim10^{5-6}$ K) component, using \\ion{O}{6} $\\lambda\\lambda1032,\\;1038$ [\\citealt{heck01} (NGC 1705)].  The existence of outflows can be inferred from the presence of absorption lines that are blueshifted with respect to the systemic velocity of the host galaxy.  We apply this technique to the warm neutral gas in a sample of 12 ULIGs using the \\ion{Na}{1} D doublet and measure individual outflow properties such as absorbing column density and outflow velocity.  \\citet{heck00} (hereafter HLSA) have used this feature to study a sample of 27 luminous infrared galaxies (LIGs; $L_{\\mathrm{IR}} > 10^{11}L_{\\sun}$) and 5 ULIGs with $z<0.12$, but generally at lower resolution ($55-170\\;\\kmsm$).  We analyze the line profiles in our targets by means of a $\\chi^2$ minimization fitting algorithm that allows for non-Gaussian intensity profiles, multiple absorbing components, and a covering fraction that is less than one.  Assuming that all absorption components with velocities greater than $\\sim70\\;\\kmsm$ are outflowing, we use a simple model of a mass-conserving free wind to compute the corresponding mass outflow rates.  We also compute star formation rates from infrared luminosities.  The ratio of the total mass outflow rate to the corresponding global star formation rate in each object is a measure of the ``reheating efficiency,'' or the efficiency with which star formation reheats the surrounding ISM and produces bubbles and outflows \\citep{mar99}.  Previous measurements of the reheating efficiencies in nearby dwarf galaxies, edge-on spirals, and lower-luminosity infrared galaxies indicates that this ratio is of order unity (\\citealt{mar99}; HLSA).  We can also track the dependence of the reheating efficiency on the mass of the galaxy, as parameterized by observed $R$- and $K^{\\prime}$-band luminosities and the kinematics of the host galaxy.  We can estimate the fraction of absorbing gas that escapes the galaxy and enters the intergalactic medium by comparing outflow velocities with escape velocities based on measured kinematics.  From the reheating efficiency and escape fraction for each object, we are able to compute the ``ejection efficiency,'' which is the ratio of the outflow rate of material that escapes into the intergalactic medium to the corresponding global star formation rate.  This last quantity is most directly related to the impact of ULIGs on the IGM.  The structure of this paper is as follows.  Section 2 describes the sample, observations, and data reduction.  In \\S 3, we present an analysis of the \\ion{Na}{1} D lines; measure kinematics, optical depths, and covering fractions; and compute absorbing column densities.  We also present equivalent width measurements of prominent absorption lines.  In \\S 4, we discuss and interpret the results; we compute mass outflow and star formation rates and compare the results to previous studies; and we discuss the ultimate fate of the outflowing gas.  In \\S 5, we summarize and conclude.  For all calculations, we assume $H_0 = 75\\; \\kmsm\\;\\mathrm{Mpc^{-1}}$ and $q_0 = 0$.  All wavelengths are vacuum wavelengths unless otherwise noted and are taken from \\citet{mort91} and/or the NIST Atomic Spectra Database\\footnote{The NIST Atomic Spectra Database is available at \\url{http://physics.nist.gov/cgi-bin/AtData/main\\_asd}.}. ",
        "conclusions": "We present the results of a moderately-high resolution spectroscopic study of a sample of 11 ULIGs.  These objects are selected from the IRAS 1 Jy sample of \\citet{ks98} as good candidates to contain massive, galaxy-scale outflows.  (This possible bias toward objects with high mass outflow rates should be considered when assessing the generality of our results.)  The goal of the study is to use prominent absorption lines to measure mass outflow rates in these galaxies.  Profiles are fit to the \\ion{Na}{1} D lines in each object to measure column densities, $N$(\\ion{Na}{1}); velocities relative to the host galaxy, $\\Delta v = v_{\\mathrm{sys}} - v$; line widths; and a covering fraction $C_f$.  Our analysis technique assumes a Gaussian velocity distribution for the absorbing gas and fits intensity profiles using a $\\chi^2$ minimization; it is useful for fitting profiles with one or two absorbing components (or more if well-separated in velocity space or otherwise constrained), profiles with a covering fraction less than unity, and even slightly saturated profiles.  We measure absorption-line components blueshifted by $\\ga 70\\;\\kmsm$ in 8 of the 11 targets ($73\\%$); this stands in between the recent measurements of HLSA and \\citet{pet01}, who find outflows in $38\\%$ and $100\\%$ of their sample, respectively.  The absorbing material in our sample have relatively low optical depths (mostly $\\tau_{1,c} < 1.0$), and the corresponding covering fractions span a wide range.  The typical maximum outflow velocity $\\Delta v_{\\mathrm{max}}$ and total outflowing \\ion{Na}{1} column density are $\\sim 300\\;\\kmsm$ and $(4-5)\\times10^{13}\\;\\mathrm{cm}^{-2}$, respectively (although one object has $\\Delta v_{\\mathrm{max}} = 1540\\;\\kmsm$).  We also measure the host galaxy kinematics in several ways and present absolute $R$- and $K^{\\prime}$-band magnitudes \\citep{vks02} in order to trace the masses and gravitational potentials of these galaxies.  From these data, we infer the following:  (1) {\\it Mass outflow rate.}  Using a simple model of a mass-conserving free wind and assuming that all absorbing components with $\\Delta v > 70\\;\\kmsm$ are outflowing, the corresponding total mass outflow rates for objects hosting a wind are in the range $13-133\\;\\smpym$.  The simplicity of our model and the difficulty of ionization corrections in the conversion from $N$(\\ion{Na}{1}) to $N$(H) likely dominate the uncertainties in $\\dot{M}$.  However, the absorbing gas is probably in harsher conditions than we assume, suggesting that our results are lower limits to the actual values.  (2) {\\it Reheating efficiency.}  The reheating efficiencies in our sample, equal to $\\dot{M}_{\\mathrm{tot}}$ divided by the corresponding global star formation rate ($\\eta \\equiv \\dot{M}_\\mathrm{tot} / \\mathrm{SFR}$), are in the range $0.1-0.7$ for galaxies hosting a wind.  These values are on average a factor of only $\\sim10$ smaller than those measured by \\citet{mar99}, who studied warm ionized and hot gas in a sample of nearby dwarf galaxies and edge-on spirals and concluded $\\eta = 0.7-5$, and \\citet{heck00} (HLSA), who studied a sample of luminous and ultraluminous infrared galaxies and concluded that $\\eta \\sim 1$.  Given that the galaxies in these studies have star formation rates that are, on average, $10-10^4$ times smaller than the SFRs in ULIGs, and that molecular gas in ULIGs may inhibit outflows, it is remarkable that their values for $\\eta$ differ by a factor of only 10.  This implies that the outflow rate of a wind scales (to within a factor of 10) with the corresponding star formation rate over a wide range of values of SFR, which in turn suggests that the physical conditions governing outflows may be somewhat similar in both ULIGs and less massive galaxies.  (3) {\\it $\\eta$ and $\\delta$ vs. host mass.}  Within our sample, the reheating efficiency apparently has no dependence on the mass of the host galaxy (as traced by emission- and absorption-line widths, central velocity dispersions, $M_R$, and $M_{K^{\\prime}}$), a result also obtained by \\citet{mar99}.  Although our subsample of objects with ejection efficiency $\\delta \\equiv \\dot{M}_\\mathrm{esc} / \\mathrm{SFR} > 0$ is small ($\\sim 5$ objects), we also find no evidence that $\\delta \\sim v_c^{-2}$, a prescription postulated by \\citet{kc98} in their theoretical study of the mass-metalliticy relation of ellipticals.  (4) {\\it Escape fraction and ejection efficiency.}  We estimate the fraction of absorbing gas in our galaxies which escapes into the surrounding medium by using host galaxy kinematics and a singular isothermal sphere potential to compute $v_\\mathrm{esc}$ and compare it to the outflowing gas velocities.  We find that 5 out of 11 objects have non-zero escape fractions, and the resulting average escape fraction of gas is $\\langle f_\\mathrm{esc} \\rangle \\equiv \\sum \\dot{M}_\\mathrm{esc}^i / \\sum \\dot{M}_{\\mathrm{tot}}^i = 0.4-0.5$.  The corresponding average ``ejection efficiency'' is $\\langle\\delta\\rangle \\equiv \\sum \\dot{M}_\\mathrm{esc}^i / \\sum \\mathrm{SFR}^i = 0.08-0.10$.  These values are fairly uncertain, and could be compromised by any sample selection bias or the small size of our sample, but are relatively insensitive to the way in which we calculate $v_\\mathrm{esc}$.  At a given $z$, the fraction of gas being expelled into the intergalactic medium that originates in ULIGs is governed by two parameters: (1) the fractional contribution of ULIGs to the total star formation rate density of the universe, and (2) the ratio of the average ejection efficiency of other star-forming galaxies to that of ULIGs, $\\langle \\delta \\rangle^{\\mathrm{sfg}} / \\langle \\delta \\rangle^{\\mathrm{ULIGs}}$.  The fractional contribution of ULIGs to star formation is modest in the local universe, but it may increase strongly with $z$ due to the observed strong increase in the number density of ULIGs with $z$.  For $\\langle \\delta \\rangle^{\\mathrm{sfg}} / \\langle \\delta \\rangle^{\\mathrm{ULIGs}} = 10$ (consistent with current observations), we find that ULIGs must host at least $50\\%$ of all star formation at a given $z$ for the fractional contribution of ULIGs to gas ejected from galaxies into the IGM to be greater than $10\\%$."
    },
    "0201/astro-ph0201237_arXiv.txt": {
        "abstract": "Recent measurements of the Cosmic Microwave Background Anisotropy have provided evidence for the presence of oscillations in the angular power spectrum. These oscillations are a wonderful confirmation of the standard cosmological scenario and allow us to derive constraints on many cosmological, astrophysical and inflationary parameters. If the discovery is confirmed by future experiments, opportunities may appear, for example, to constrain dark energy, variations in fundamental constants and neutrino physics. ",
        "introduction": "In the last $2$ years important progress has been made in the study of the Cosmic Microwave Background (CMB) Anisotropies.  With the TOCO$-97/98$ (\\cite{torbet},\\cite{miller}) and Boomerang-$97$ (\\cite{mauskopf}) experiments a firm detection of the first peak in the CMB anisotropy angular power spectrum has been obtained. The presence of this peak is generally expected in models of structure formation with a nearly-scale invariant spectrum of primordial perturbations like the one produced after inflation. In the framework of adiabatic Cold Dark Matter (CDM) models, the position, amplitude and width of this peak provide strong supporting evidence for the inflationary predictions of a low curvature (flat) universe and a scale-invariant primordial spectrum (\\cite{knox}, \\cite{melchiorri}, \\cite{tegb97}).  A first analysis of a small fraction of data from the BOOMERANG $1998/1999$ Long Duration Ballooning (BOOM/LDB) campaign (\\cite{debe00}, \\cite{lange}) and of observations from the MAXIMA experiment (\\cite{hanany}, \\cite{balbi}) further confirmed the presence of this feature at high significance. However, the finding of a suppressed second peak in the CMBR anisotropy resulted in a rather large value for the baryon density, $\\Omega_b h^2 = 0.032^{+0.005}_{-0.004}$ at $68\\%$ CL~\\cite{th1}, while the experimental data on primordial $^4He$ and $D$ abundances, prefer smaller values, $\\Omega_b h^2 = 0.020 {\\pm} 0.002$ (\\cite{burles}) (see also \\cite{avelino},\\cite{Lisi}). Many authors addressed the issue of this tension between the determination of $\\Omega_b h^2$ from CMBR data and Standard Big Bang Nucleosynthesis (SBBN) \\cite{th4,peloso,th7,Hansen,Kaplin,dibari}.  The new experimental data from BOOMERANG (\\cite{netterfield}) and DASI (\\cite{halverson}) have refined the data at larger multipole and now clearly suggest the presence of a second peak in the spectrum and a smaller value for the baryonic fraction, in agreement with SBBN. Moreover, this result confirms the model prediction of acoustic oscillations in the primeval plasma and shed new light on various cosmological and inflationary parameters (\\cite{debe01}, \\cite{wang}, \\cite{pryke}). The new results from MAXIMA (\\cite{lee}, \\cite{stompor}) are of lower precision, but are consistent with both DASI and BOOMERANG.  This paper is organized as follows: Section II is a brief introduction about why we expect oscillations in the CMB spectrum. In Section III I will discuss the statistical significance of the peaks measured by BOOMERANG, DASI and MAXIMA and their location and amplitude. In section IV I will review the implications for the cosmological parameters in the framework of the standard CDM model of structure formation. In section V I will discuss some non-standard aspect of parameter extraction. Finally, in section VI, I will give my conclusions. ",
        "conclusions": "The recent CMB data represent a beautiful success for the standard cosmological model. The acoustic oscillations in the CMB angular power spectrum, a major prediction of the model, have now been detected at $\\sim 5 \\sigma$ C.L. for the first peak and $\\sim 2 \\sigma$ C.L. for the second and third peak. Furthermore, when constraints on cosmological parameters are derived under the assumption of adiabatic primordial perturbations the following results are obtained:  \\begin{itemize}  \\item The curvature of the universe is zero, i.e. the universe is flat, in agreement with the predictions of the theory.  \\item The power spectrum of the primordial perturbations is nearly scale-invariant, again a prediction of the model.  \\item The amount of density in baryons is in agreement with independent observations of primordial abundances and standard big-bang nucleosynthesis.  \\item The optical depth is constrained to be $\\tau_c <0.3$, and the universe recombined, in agreement with the overall scenario.  \\item Some form of non baryonic dark matter must be present, as requested by a large set of independent observations.  \\item The age of the universe is consistent with at least $3$ independent constraints.  \\item When information from complementary datasets, like constraint on $h$ or from large scale structure, are included in the analysis, the CMB data suggest a presence for a cosmological constant in agreement with the SN-Ia result.  \\end{itemize}  All these results strongly suggest that the inflationary scenario of structure formation is coherent in its simplest form. In few words, the hoped-for miracle.  As we saw in the previous section, modifications to this model, like adding isocurvature modes or topological defects, are in agreement with the observations, but are not required by the data and are reasonably constrained when complementary datasets are included in the analysis.  Since the model is in agreement with the data and all the most relevant parameters are starting to be constrained within a few percent accuracy, the CMB is becoming a wonderful laboratory for investigating the possibilities of new physics. With the promise of large data sets from Map, Planck and SNAP satellites, opportunities may be open, for example, to constrain dark energy models, variations in fundamental constants and neutrino physics."
    },
    "0201/astro-ph0201551_arXiv.txt": {
        "abstract": "We present new Hubble Space Telescope images of the gravitational lens \\lens, which allow us to characterize the lens galaxy and update the determination of the Hubble constant ($H_0$) from this system.  The $I$-band image shows that the lens galaxy is a face-on spiral galaxy with clearly delineated spiral arms.  The southwestern image of the background quasar passes through one of the spiral arms, explaining the previous detections of large quantities of molecular gas and dust in front of this image.  The lens galaxy photometry is consistent with the Tully-Fisher relation, suggesting the lens galaxy is a typical spiral galaxy for its redshift.  The lens galaxy position, which was the main source of uncertainty in previous attempts to determine $H_0$, is now known precisely.  Given the current time delay measurement and assuming the lens galaxy has an isothermal mass distribution, we compute $H_0=44\\pm 9$~km~s$^{-1}$~Mpc$^{-1}$ for an $\\Omega_m =0.3$ flat cosmological model.  We describe some possible systematic errors and how to reduce them.  We also discuss the possibility raised by Courbin et al.\\ (2002), that what we have identified as a single lens galaxy is actually a foreground star and two separate galaxies. ",
        "introduction": "Although the Hubble Space Telescope (HST) Key Project to Measure the Hubble Constant has been successfully completed \\citep{freedman01}, it is important to pursue completely different methods of determining the Hubble constant ($H_0$).  Independent methods provide a consistency check, and may eventually surpass the 10\\% accuracy of the local distance-scale methods employed by the Key Project.  It is also important to measure the expansion rate directly at cosmological redshifts, in case the Galaxy is in a locally underdense or overdense region \\cite[see, e.g.,][]{wu95}.  These goals will grow in importance in the coming years due to the degeneracies in analyses of cosmic microwave background anisotropies between many cosmological parameters and $H_0$ \\cite[e.g.,][]{bond94,eisenstein98}.  A promising approach to determining $H_0$ that is independent of the local distance scale uses gravitational lens time delays (\\citealt{refsdal64}; for recent summaries see, e.g., \\citealt{schechter00}, \\citealt{koopmans99}).  This method is ultimately limited by systematic uncertainties in the mass models of the lens galaxies, but a necessary first step is to obtain the basic observational constraints---astrometry, photometry, and redshifts of the lens and source---for the systems with measured time delays.  For the gravitational lens \\lens, this has been especially challenging because the system is near the Galactic plane ($b=-5\\fdg 7$).  As a result, most of the information has come from radio and infrared observations.  \\citet{rao88} suggested it was a lens due to its radio morphology. \\citet{subrahmanyan90} and \\citet{jauncey91} provided more conclusive evidence; the system has two bright point sources (NE and SW) embedded in a fainter Einstein ring.  Strong molecular absorption features were detected at $z_l=0.886$ \\citep{wiklind96,gerin97}, largely in front of SW \\citep{frye97,swift01}, and presumably due to the lens galaxy.  \\citet{lovell96} found $z_a=0.19$ \\ion{H}{1} absorption in front of NE, of unknown provenance.  \\citet{lovell98} measured a time delay of $26^{+4}_{-5}$~days between NE and SW at 8.6~GHz.  By deconvolving ground-based infrared images, \\citet{courbin98} detected the quasar images and found hints of the lens galaxy.  \\citet{lidman99} determined the quasar redshift of $z_s=2.507$ by infrared spectroscopy.  From the presence and structure of the molecular absorption system, \\citet{wiklind98} argued that the lens galaxy is probably a spiral galaxy seen nearly face-on, which is consistent with the large differential extinction between NE and SW ($\\Delta E_{B-V} \\approx 3$) observed by \\citet{falco99}.  However, although \\citet{lehar00} detected the lens galaxy in $H$-band and shallow $I$-band HST images, they could not determine its position or structure accurately.  These authors found that the uncertainty in the lens galaxy position dominated the uncertainty in the value of $H_0$ inferred from the time delay.  In this paper we present new HST optical images that confirm the lens galaxy is a face-on spiral galaxy, and allow its position and optical magnitudes to be measured accurately.  These data are discussed in \\S~\\ref{sec:lensgalaxy}.  In \\S~\\ref{sec:tullyfisher}, we place the lens galaxy on the Tully-Fisher relation using the HST photometry and two different estimators of the galaxy mass---one from the lens geometry, and one from the velocity shift measured by \\citet{wiklind98}.  In \\S~\\ref{sec:models}, we incorporate our measurement of the position of the lens galaxy into a simple lens model to arrive at an updated estimate of $H_0$.  We also discuss systematic errors and compare our model with previous models.  After this work was completed, we learned that \\citet{courbin02} independently analyzed the same HST data, along with new and archival near-IR data.  They argue that what we have identified as the bulge of the lens galaxy is actually a foreground star, and that there is a second lens galaxy.  In the final section of this paper, we describe the strengths, weaknesses, and future tests of these competing hypotheses, and discuss future prospects for reducing systematic errors in $H_0$ from this particular gravitational lens. ",
        "conclusions": "\\label{sec:summary}  The new HST $I$-band image confirms that the lens galaxy of \\lens~is a nearly face-on spiral galaxy.  Assuming that the compact component near the center of this galaxy is the galaxy bulge, we have accurately measured the galaxy position relative to the quasar images, thereby completing the basic data for this system.  Furthermore, we have computed $H_0$ given the current best measurement of the time delay, assuming that there is only one lens galaxy, and further assuming that the lens galaxy has a flat rotation curve and a massive dark halo, as appears to be the case for nearby massive spiral galaxies.  The resulting $H_0$ is lower than the widely accepted value obtained by the Key Project.  We have described some possible sources of bias, but to force agreement with the Key Project seems to require an unnatural mass model for the lens galaxy.  After this work was done, we learned that \\citet{courbin02} independently analyzed the same HST data as presented in \\S~\\ref{sec:lensgalaxy} (along with the $IHK$ data of \\cite{lehar00} and a new Gemini $K$-band image).  The main difference between their treatment of the $VI$ data and ours is that they deconvolved the images before interpreting them.  Their interpretation differs from ours in two important respects:  \\begin{enumerate}  \\item \\citet{courbin02} conclude object P is a foreground star rather than the bulge of the face-on spiral galaxy.  In support of this claim, P is unresolved and its position in a color-magnitude diagram is consistent with other bulge dwarf stars for this field.  Based on the mean density of field stars that were detected in the $I$-band image ($\\sim 0.8$~arcsec$^{-2}$), there is a 10\\% chance for a randomly placed star to lie within $0\\farcs2$ of the center of the galaxy.  The photometry is also consistent with a spiral galaxy with P as its bulge (see \\S~\\ref{sec:tullyfisher}), so both interpretations are reasonable.  \\item \\citet{courbin02} conclude there is a second lens galaxy, based on faint and extended flux between P and the SW quasar that is detected only in the $H$-band image of \\citet{lehar00}.  Its position is rather uncertain (with a quoted error of 80~mas) but is consistent with the position of radio component E.  In this scenario, the $H$-band flux and component E are due to a second deflector that is radio-loud.  In our scenario, these features (if they are real, and not associated with the quasar) could be due to star formation in the western spiral arm of lens galaxy.  \\end{enumerate}  In short, our interpretation has the virtue of simplicity.  The interpretation of \\citet{courbin02} is complicated but has the merit of explaining a few puzzling features of the data: a possible offset between P and the center of the spiral arms, the diffuse $H$-band flux, and radio component E.  It is not clear to us, from the present data, how seriously these features should be taken.  If P is indeed a star, the lens galaxy position we report in this paper is wrong.  Adopting the position of \\citet{courbin02} and the SIE model described in \\S~\\ref{subsec:update}, the inferred Hubble constant would rise to $H_0= 107\\pm 30$~km~s$^{-1}$.  If, however, there is a second galaxy between the quasar images, then it may be impossible to determine $H_0$ from this system because the galaxies would be difficult to characterize with current telescopes, and the mass model would be complex.  In particular it is not possible to state generally (without detailed modeling) whether a two-deflector model would alleviate or exacerbate the incompatibility of the Key Project value of $H_0$ and the time delay measured for this system.  Therefore, for the purpose of determining $H_0$ with this system, the highest priority should be establishing which of these competing hypotheses are correct.  It would be difficult to establish the identity of P via spectroscopy, because P is faint and the field is crowded.  But, as \\citet{courbin02} pointed out, it would be possible to prove P is a star by measuring its proper motion, which should be $\\sim 4$~mas~yr$^{-1}$ for a bulge star (for which the solar reflex motion of $\\sim 200$~km~s$^{-1}$ is dominant).  A negative result would not be conclusive.  The reality of the second deflector might be tested by seeking to detect both deflectors in a single image---perhaps a considerably deeper $I$-band image with the HST's new Advanced Camera for Surveys.  Assuming for the moment that our interpretation of the HST images is correct, we suggest the following steps to reduce the systematic errors and sharpen the determination of $H_0$ from this system.  First, the uncertainty in the time delay should be reduced from 20\\% to 3\\% or lower, so that it makes no significant contribution to the overall uncertainty.  \\citet{lovell98} noted that their full light curves were consistent with a broad range of time delays, ranging from 12 to 30 days.  Only when the data were restricted to a subset containing a single ``bump'' did the distribution of possible time delays become approximately Gaussian.  A time delay based on multiple features in a full set of light curves would be more secure.  Second, the redshift of galaxy G2 ($I=20.7$) needs to be measured spectroscopically to test whether it is the source of the $z_a=0.19$ \\ion{H}{1} absorption. The plausibility of the association with G2 could also be checked by mapping the \\ion{H}{1} absorption with a higher signal-to-noise ratio than the map of \\citet{lovell96}, in order to search for absorption in front of SW, and to measure any velocity shift.  Third, and most important, the degeneracies of lens models must be broken by making use of more observational constraints than the basic data presented in Table~\\ref{tbl:data}.  For example, we did not investigate realistic models for spiral galaxies that include contributions from the bulge, disk, and halo \\cite[see, e.g.,][]{keeton98,koopmans98}.  More generally, many authors have argued that unless a wide class of parameterized models are considered \\cite[e.g.,][]{kochanek91,bernstein99}, or even a non-parametric model \\cite[e.g.,][]{williams00}, the uncertainty in $H_0$ will be underestimated.  To investigate a broader range of models, more constraints are obviously required.  For the simple models described in \\S~\\ref{sec:models}, we have ignored two possible sources of such constraints: the milliarcsecond-scale radio morphologies of the quasar cores, and the Einstein ring.  The quasar cores have been mapped with very long baseline interferometry (VLBI) at radio frequencies ranging from 843~MHz to 43~GHz.  Maps with angular resolution $\\gtorder 10$~mas show the expected parity-reversed substructure in NE and SW \\cite[see, e.g.,][] {jauncey91,patnaik96}, but maps with higher angular resolution are difficult to interpret.  Component SW, at least, appears to be scatter-broadened by plasma in the lens galaxy \\citep{jones96,guirado99}.  It has been claimed that both NE and SW exhibit intrinsic variability in source structure \\citep{jin99}. Furthermore, NE has a $\\sim 10$~mas linear jet with no obvious counterpart in SW \\citep{garrett96,guirado99}, which may indicate the presence of small-scale substructure in the lens.  It is hard to see how the existing observations can be dramatically improved with current telescopes, but a modeling effort may help to determine whether mass subtructure is required, and if so, whether the mass scale of the perturbation is large enough to affect the time delay.  The geometry of the Einstein ring is probably the best hope for breaking the remaining degeneracies of lens models.  The situation is analogous to the case at optical wavelengths, where the stretched images of quasar host galaxies have been used to break some degeneracies in mass models of the time delay lens PG~1115+080 (Kochanek, Keeton, \\& McLeod 2001; although see Saha \\& Williams 2001).  We have made some preliminary attempts to use the existing VLA maps for this purpose, by employing the curve-fitting algorithm of \\citet{keeton01a}, but with the angular resolution of the present data it is difficult to measure the ring with the needed accuracy.  There is room for improvement upon existing radio maps of the ring, which tend to be short snapshots because the radio cores are so bright ($\\sim 3$ Jy).  (In most of the VLBI observations, the ring is resolved out.)  These snapshots are limited in dynamic range by poor sampling in Fourier space rather than thermal noise.  Major improvements would follow from deep radio observations with high angular resolution ($<0\\farcs1$) and high dynamic range ($>10^5$). Methods for the precise interpretation of high-dynamic-range interferometric observations of lensed sources are well developed and have been shown to discriminate between different mass models \\cite[see, e.g.,][]{kochanek92,kochanek95,chen95,ellithorpe96}."
    },
    "0201/astro-ph0201417_arXiv.txt": {
        "abstract": "This paper describes the acquisition and analysis of imaging data for the identification of galaxies associated with $z \\sim 4$ damped \\lya systems.  We present deep $BRI$ images of three fields known to contain four $z \\sim 4$ damped systems.  We discuss the reduction and calibration of the data, detail the color criteria used to identify $z \\sim 4$ galaxies, and present a photometric redshift analysis to complement the color selection.  We have found no galaxy candidates closer to the QSO than $7''$ which could be responsible for the damped \\lya systems.  Assuming that at least one of the galaxies is not directly beneath the QSO, we set an upper limit on this damped \\lya system of $L < L_{LBG}^*/4$.  Finally, we have established a web site to release these imaging data to the public (http://kingpin.ucsd.edu/$\\sim$dlaimg). ",
        "introduction": "In this era of large telescopes, it is almost routine to identify galaxies at high redshift.  Deep narrow band imaging \\citep[e.g.][]{hu98} and the Lyman break technique \\citep{steidel93}, among other approaches, have led to the discovery of over 1000 $z > 2$ galaxies where just 10 years ago only a handful were known. Unfortunately, because these techniques are magnitude limited, they are primarily sensitive to the high luminosity tail of the protogalactic distribution.  For example, the Lyman break galaxies exhibit a two-point correlation function which suggests they are massive galaxies associated with significant large-scale structures \\citep{adlb98} that will evolve into the most massive galaxies today.  To study the formation of typical galaxies, therefore, one must circumvent this luminosity selection.  Even before the recent outburst of high redshift galaxy identifications, researchers had already discovered a significant population of high $z$ protogalaxies in absorption: the damped \\lya systems \\citep[e.g.][]{wol86}. Because the selection of these galaxies is based on H~I gas cross-section, they are predicted to sample the bulk of the galactic mass distribution \\citep{kau96} and therefore provide a broader picture of galaxy formation in the early universe. Attempts to identify in emission the galaxies which give rise to damped systems, however, have met with limited success.  Djorgovski and collaborators \\citep{djg96,djg97} have made a few discoveries, but their success rate is uncertain. More recently, \\cite{warren01} have analysed deep HST images of fields around quasars known to exhibit damped \\lya systems and have identified small, relatively faint galaxies close to the quasar claimed to be the damped \\lya systems. In both of these cases, the surveys have focused on galactic candidates within a few arcseconds of the background quasar. Although this is an important aspect of damped \\lya research, there are unique advantages to extending the galaxy surveys to a much larger area. In particular, one can estimate the mass of the damped systems -- and therefore the bulk of the high $z$ protogalactic population -- by measuring their clustering properties \\citep{gawiser01a}.  This paper describes the observational and analysis techniques we have implemented to select high redshift galaxy candidates in deep $BRI$ images of three fields containing four known damped \\lya systems. It marks the first step in an observing program designed to search in emission for the galaxies giving rise to or clustered with damped \\lya systems at high redshift. The present observations focus on fields with known $\\zabs \\sim 4$ damped \\lya systems, i.e.\\ the galaxies associated with these systems will be $B$-band dropouts. While follow-up observations of $z \\sim 4$ galaxies are considerably more difficult than at $z \\sim 3$ \\citep{steidel99}, we have been limited by the absence of a $U$-band imager at the Keck Observatories. Table~\\ref{tab:dla} lists the known properties of the four damped \\lya systems from the three quasar fields described in this paper.  Each system exhibits a neutral hydrogen column density which exceeds the statistical survey threshold $(\\N{HI} = 2 \\sci{20} \\cm{-2}$) established by \\cite{wol86}. For two of the quasars we have obtained high resolution spectroscopy \\citep{pro99,pro00,pro01} with HIRES \\citep{vogt94} on the Keck~I telescope and have measured their [Fe/H] metallicity. The damped system toward {\\qof} is remarkable for exhibiting by far the highest metallicity of any $z>3$ damped \\lya system to date while the $z = 4.20$ system toward {\\qon} has among the lowest metallicity.  \\begin{table}[ht] \\begin{center} \\caption{{\\sc DAMPED LYA SYSTEMS \\label{tab:dla}}} \\begin{tabular}{lcccc} \\tableline \\tableline Quasar & $z_{QSO}$ & $z_{DLA}$ & $\\N{HI}$\\tablenotemark{a} & [Fe/H] \\\\ \\tableline PSS0132+1341  & 4.15 & 3.93 & 20.3 & \\\\ BR0951$-$0450 & 4.37 & 4.20 & 20.4 & $< -2.6$\\tablenotemark{b} \\\\ &      & 3.86 & 20.6 & $-2.1$\\tablenotemark{b} \\\\ PSS1443+2724  & 4.41 & 4.22 & 20.8 & $-1.0$\\tablenotemark{c} \\\\ \\tableline \\end{tabular} \\tablenotetext{a}{\\cite{storr00}} \\tablenotetext{b}{\\cite{pro99}} \\tablenotetext{c}{\\cite{pro00,pro01}} \\end{center} \\end{table}   First results on one of these fields (\\qon) have been presented elsewhere \\citep{gawiser01a} and a companion paper \\citep[][Paper II]{gawiser01b} describes the reduction and analysis of the multi-slit spectroscopy of all three fields. In $\\S$~\\ref{sec-obsred} we discuss the imaging observations and reduction process, the photometric calibration is described in $\\S$~\\ref{sec-photcalib}, we present the photometry and astrometry for all detected objects in $\\S$~\\ref{sec:photom}, $\\S$~\\ref{sec:results} presents the results, and $\\S$~\\ref{sec:summary} gives a brief summary. Unless otherwise noted, all magnitudes have been placed on the $AB_{95}$ scale by offsetting Johnson-Cousins magnitudes: $B_{AB} = B_{JC} - 0.11; R_{AB} = R_{JC} + 0.20; I_{AB} = I_{JC} + 0.45$ \\citep{fuk96}.  \\begin{table*} \\begin{center} \\caption{{\\sc JOURNAL OF OBSERVATIONS \\label{tab:obs}}} \\begin{tabular}{lcccccc} \\tableline \\tableline QSO Field & UT & Filter & Eff Exp & FWHM & SB$^{1\\sigma}$ & SB$^{3\\sigma}$ \\\\ & & & (s) & ($''$) & (Mag$_{AB}$/$\\square ''$) & (Mag$_{AB}$/FWHM$^2$) \\\\ \\tableline PSS0132+1341  & Jan 13 1999 & B & 3750 & 1.05 & 29.5 & 28.3 \\\\ &             & R & 1680 & 0.85 & 28.8 & 27.8 \\\\ &             & I & 1080 & 0.58 & 28.0 & 27.4 \\\\ BR0951$-$0450 & Jan 13 1999 & B & 4500 & 0.85 & 29.8 & 28.8 \\\\ &             & R & 1620 & 0.75 & 29.0 & 28.1 \\\\ &             & I & 1560 & 0.63 & 28.1 & 27.4 \\\\ PSS1443+2724  & Mar 23 1999 & B & 3600 & 1.08 & 29.7 & 28.4 \\\\ &             & R & 1200 & 0.80 & 28.8 & 27.9 \\\\ &             & I & 750  & 1.02 & 27.8 & 26.6 \\\\ \\tableline \\end{tabular} \\end{center} \\end{table*} ",
        "conclusions": "\\label{sec:summary}  We have presented an analysis of $BRI$ images for three fields containing quasars with four known damped \\lya systems at $z \\sim 4$.  The data was reduced with standard procedures and calibrated with Landolt standard stars taken close to the time of the observations. We measured $BRI$ magnitudes for all of the objects in each field using isophotal apertures and measured $B-R$ and $R-I$ colors for those galaxies detected in the $R$-images.  A rough number count analysis agrees well with other studies in the literature to similar magnitude limits.  By searching for $B$-band dropouts and measuring photometric redshifts, we have culled a list of candidates at $z \\sim 4$ throughout each field for follow-up spectroscopy.  We identified $\\approx 30 - 40$ candidates per field to a limiting magnitude of $I_{AB} < 25.5$.  For the three fields there are only two dropouts detected within $10''$ of the quasar, each at an impact parameter of $\\approx 7''$. Although these two galaxies could give rise to one or two of the damped \\lya systems, we expect it is more likely that they are either clustered with the damped system or at a different redshift altogether.  Therefore, we have no convincing detections of the galaxy in emission for four damped systems in the three fields. Either the galaxies are fainter than our detection limit ($\\approx L_{LBG}^*/4$) or located under the PSF of the background quasar.  Nevertheless, our results enable a study of the large-scale clustering of galaxies associated with the damped \\lya systems.  Indeed, this is the focus of the analyses presented in companion papers (Paper~II)."
    },
    "0201/astro-ph0201403_arXiv.txt": {
        "abstract": "Over the last few years X-ray observations of broad-line radio galaxies (BLRGs) by \\textit{ASCA}, \\textit{RXTE} and \\textit{BeppoSAX} have shown that these objects seem to exhibit weaker X-ray reflection features (such as the iron K$\\alpha$ line) than radio-quiet Seyferts. This has lead to speculation that the optically-thick accretion disc in radio-loud active galactic nuclei (AGN) may be truncated to an optically-thin flow in the inner regions of the source. Here, we propose that the weak reflection features are a result of reprocessing in an ionized accretion disc. This would alleviate the need for a change in accretion geometry in these sources. Calculations of reflection spectra from an ionized disc for situations expected in radio-loud AGN (high accretion rate, moderate-to-high black hole mass) predict weak reprocessing features. This idea was tested by fitting the \\textit{ASCA} spectrum of the bright BLRG \\threec\\ with the constant density ionized disc models of Ross \\& Fabian. A good fit was found with an ionization parameter of $\\xi \\sim 4000$~erg~cm~s$^{-1}$ and the reflection fraction fixed at unity. If observations of BLRGs by \\textit{XMM-Newton} show evidence for ionized reflection then this would support the idea that a high accretion rate is likely required to launch powerful radio jets. ",
        "introduction": "\\label{sect:intro}  In the standard disc-corona paradigm for the X-ray emission from active galactic nuclei (AGN), the hard X-ray source is energized within a magnetically dominated corona and illuminates the optically-thick accretion disc below it (see \\citealt{bel99} for a recent review). The reprocessed emission from the disc imprints spectral features on the observed spectrum, most notably the \\fe\\ line (which may be broadened due to relativistic effects; \\citealt{fab89,tan95}; \\citealt{fab00}) and a Compton reflection hump at about 20--30~\\kev\\ due to the down scattering of higher energy photons \\citep[e.g.,][]{pou90}. By measuring the strength of these reflection features and comparing to models of X-ray reflection, it is possible, in principle, to constrain the solid-angle subtended by the reflector (as seen from the hard X-ray source), and hence the geometry of the accretion flow a few Schwarzschild radii away from the black hole. Thus, measurement of the X-ray reflection signatures is not only a valuable test for the disc-corona model, but also an important method to probe the central engines of AGN.  Over the last few years observations of radio-loud AGN, particularly broad-line radio galaxies (BLRGs), by \\textit{ASCA} and \\textit{BeppoSAX} have indicated that these sources may have weaker hard X-ray reprocessing features than radio-quiet AGN at similar luminosities \\citep*[e.g.,][]{eh98,gr99,esm00,gr01}. This has lead some authors to speculate that the thin, optically-thick accretion disc in radio-loud AGN changes to a hot, optically-thin flow before reaching the black hole. Similar conclusions have also been made about Galactic black hole candidates in their low/hard state: they also appear to exhibit weaker reflection features than expected if the standard accretion disc extends down to the last stable orbit \\citep[e.g.,][]{zds97,zds98,zds99,dz99}. Moreover, strong radio emission from black hole candidates tends to be seen only when they are in their low/hard state \\citep{fen01}. A truncated disc in radio-loud objects has also been supported by the recent theoretical work of \\citet{mei01} who considered the strength of the poloidal magnetic field component expected in a standard \\citet{ss73} disc as compared to one in a ADAF-like flow \\citep[e.g.,][]{ny95}. He concluded that only the combination of a rapidly spinning black hole and a geometrically thick accretion flow (i.e., an ADAF) could produce the required jet power for a radio-loud AGN.  However, the problem of accurately determining the amount of reflection in an X-ray spectrum depends on the sophistication of the reflection models available to fit the data. Over the past 10 years numerous calculations of X-ray reflection from Compton thick matter have been performed to compare against the data \\citep*[e.g.,][]{ros93,zy94,mz95,ros99,nkk00,brf01}. These computations have increased in sophistication over the years: first relaxing the assumption of neutral material and then enforcing the condition of hydrostatic balance on the illuminated material.  Ionization effects can have an important impact on the shape of the reflection spectrum and on the features imprinted on it \\citep{ros99}. In fact, if the disc surface is highly ionized the reflection features become very weak, and if not taken into account, \\emph{a low reflection fraction can be measured}. Such ionized disc models have been shown to account for the measured low reflection fractions in Cyg~X-1 \\citep{you01} and the black-hole transient Nova Muscae \\citep{dn01}. If this is the case, there is no need for a truncated optically-thick accretion disc in these sources and any jet emission would then arise from a standard disc that extends down to the last stable orbit. However, the inner parts of the disc would be highly ionized, likely due to extreme irradiation from the corona, and possibly radiation-pressure dominated.  In this Letter, we investigate the idea that the low reflection fractions measured in the hard X-ray spectra of radio-loud AGN are simply the result of reprocessing from an ionized accretion disc. In the next Section we calculate the reflection spectra from ionized discs which may be appropriate for radio-loud AGN, and show that they will exhibit weak reflection features. In Section~\\ref{sect:res} we apply ionized disc models to the \\textit{ASCA} spectrum of the BLRG \\threec. Finally, we conclude by arguing that, in light of recent results on the radio power of active galaxies, ionized accretion discs should exist in most broad-line radio galaxies. ",
        "conclusions": "\\label{sect:discuss}  In this Letter we have presented the idea that the weak reprocessing features seen in the \\textit{ASCA} spectra of many radio-loud AGN are the result of reflection off an ionized disc. With this explanation there does not need to be a significant difference in the accretion geometry between radio-loud and radio-quiet sources. This conclusion is supported by evidence that the traditional bimodal distribution in radio power of AGN is no longer valid. \\citet{llr01} used results from the FIRST Bright Quasar Survey to fill in the gap between the radio-quiet and radio-loud populations, and concluded there was a continuous distribution of radio-power that is correlated with both black hole mass and accretion rate. Similarly, \\citet{hp01} employed high resolution radio and optical imaging of Seyfert galaxies to show that many of these objects that were thought to be radio-quiet are actually radio-loud (using the $R \\equiv L_{\\nu}\\mathrm{(6\\ cm)}/L_{\\nu}\\mathrm{(B)} > 10$ definition; \\citealt{vis92}). These authors also argued that there is a radio power-optical luminosity correlation in active galaxies that stretches from Seyfert galaxies up to luminous quasars.  There still remains the problem that only a small percentage of AGN exhibit powerful, kpc-scale radio jets. The results of \\citet{llr01} suggest that a combination of high accretion rate and large ($> 10^{8.5}$~M$_{\\odot}$) black hole mass is required to launch ultra-relativistic jets. The radio luminosity-optical luminosity correlations\\footnote{The correlations seem to exist for both radio-loud and radio-quiet objects, but with the slope being steeper for the radio-loud ones.} discussed by \\citet{hp01} also point toward an accretion rate dependence. The actual physical mechanism responsible for the radio emission and jets presumably is a magnetohydrodynamic coupling between the accretion flow and a spinning black hole \\citep[e.g.,][]{bz77,mt82,ree82,bbr84,li00a,li00b}, but the details are far from worked out \\citep[e.g.,][]{ga97,lop99,mei99,mei01}. However, the key parameters must be the black hole mass and accretion rate because it is likely that rapidly spinning black holes are found in many AGN, irrespective of their radio power (e.g., MCG--6-30-15, \\citealt{iwa96,iwa99,wil01}). Indeed, studies of the hard X-ray background suggest that most supermassive black holes are rapidly spinning \\citep{fi99,erz02}. Perhaps only rapid accretion onto a very massive spinning black hole can provide the necessary energy to launch kpc-scale radio jets.  However, observational evidence suggests that at least one more parameter other than the black hole mass and the accretion rate is needed to explain the triggering of powerful jets. Recent \\textit{XMM-Newton} observations of some high-luminosity radio-quiet Seyferts have shown evidence for ionized \\fe\\ lines \\citep{ree01,pou01,orr01}, indicating that the accretion disc is becoming more ionized as the luminosity and, likely, the accretion rate increases \\citep{br02}. A widening zone of extreme ionization on the accretion disc would also explain the \\textit{ASCA} observations of radio-quiet quasars which found that the EW of the \\fe\\ line becomes undetectable at a 2--10~\\kev\\ luminosity of 10$^{46}$~erg~s$^{-1}$ (the so-called X-ray Baldwin effect; \\citealt{it93,nan97,rt00}). Although most of these quasars have an accretion rate close to Eddington and black hole masses $> 10^{8}$~M$_{\\odot}$ \\citep[e.g.,][]{llr01}, they do not possess powerful radio jets. Therefore, additional quantities, such as the magnetic field strength and configuration close to the black hole, must also be important in determining the strength of the radio jets.  Nevertheless, it is possible that the weak X-ray reprocessing features from radio-loud sources provide corroborating evidence that a high accretion rate is important for the production of powerful jets. Ultrasoft Seyferts (the subset of narrow-line Seyfert~1 galaxies with unusual X-ray properties; \\citealt{bra99,lei99a,lei99b,vau99b}) which are likely accreting close to their Eddington limit, have shown evidence for ionized accretion discs in their X-ray spectra \\citep{com98,com01,tgn98,vau99a,bif01,tur01a,tur01b}. These objects would not exhibit large scale radio jets because they do not have massive enough black holes, or, perhaps, massive enough bulges to contain an extensive reservoir of gas. Therefore, if the weak X-ray reflection features in radio-loud AGN are due to ionization effects, then it is likely to be a result of their high accretion rate, supporting the view that a large accretion rate is required for the production of relativistic jets. High quality spectral observations of radio-loud AGN by \\textit{XMM-Newton} will be able to test this hypothesis."
    },
    "0201/astro-ph0201129_arXiv.txt": {
        "abstract": " ",
        "introduction": "Explaining the characteristic period of the quasicyclic activity oscillations of stars with the Sun included is one of the challenges for stellar physics. The main period is an essential property of the dynamo mechanism. Solar dynamo theory is reviewed here in  the special context of the cycle-time problem. The parameters of the convection zone turbulence do not easily provide us with the 22-year time scale for the solar dynamo. Even for the boundary layer dynamos, it is  only possible if a `dilution factor' in the turbulent electromotive force smaller than unity  is  introduced which parameterizes the intermittent character of the MHD-turbulence.  A notable number of interesting phenomena have been investigated in the search for the solution of this problem: flux-tube dynamics, magnetic quenching, parity breaking, and chaos. Nevertheless, even the simplest observation -- the solar cycle period of 22 years -- is hard to explain (cf. DeLuca and Gilman, 1991; Stix, 1991; Gilman, 1992; Levy, 1992; Schmitt, 1993; Brandenburg, 1994a; Weiss, 1994). How can we understand the existence of the large ratio of the mean cycle period and the correlation time of the turbulence? Three  main observations are basic in this respect: \\vspace*{-0.3truecm} \\begin{itemize} \\parskip0pt \\parsep0pt \\itemsep0pt \\item There is a factor of about 300 between the solar cycle time and the Sun's rotation period. \\item This finding is confirmed by stellar observations (Fig.~\\ref{cycles}). \\item The convective turnover time near the base of the convection zone is very similar to the solar rotation period. \\end{itemize} \\vspace*{-0.3truecm} The problem of the large {\\it observed\\/} ratio of cycle and correlation times, \\begin{equation} \\frac{\\tau_{\\rm cyc}}{\\tau_{\\rm corr}} \\gsim 10^2, \\label{cyc} \\end{equation} constitutes the primary concern of dynamo models. In a thick convection shell this number reflects (the square of) the ratio of the stellar radius to the correlation length and numbers of the order 100 in (\\ref{cyc}) are possible. For the thin boundary layer dynamo, however, the problem becomes more dramatic and is in need of an extra hypothesis.  The activity period of the Sun varies strikingly about its average from one cycle to another. Only a nonlinear theory will be able to explain the non-sinusoidal (chaotic or not) character of the activity cycle (Fig.~\\ref{quality}). A linear theory is only concerned with the {\\it mean} value of the oscillation frequency. \\begin{figure} \\psfig{figure=stellcyc.ps,width=12.5cm,height=6.5cm,bbllx=-50pt,bblly=360pt,bburx=558pt,bbury=720pt} \\caption{Stellar cycles: the ratio of the cycle time to the rotation period after Saar and Brandenburg (1999). The solid line fits the observations for all stars and leads to the weak proportionality $\\omega_{\\rm cyc} \\propto \\Omega^{-0.1}$ while a previous fit (dashed line) provided $\\omega_{\\rm cyc} \\propto \\Omega^{-0.2}$ (Brandenburg {\\it et al.}, 1998).} \\label{cycles} \\end{figure} \\begin{figure} \\mbox{\\psfig{figure=quality.ps,angle=270,height=6truecm,width=7.5truecm,bbllx=12pt,bblly=62pt,bburx=576pt,bbury=757pt} \\hfill \\psfig{figure=frick.ps,height=5.5truecm,width=7.5truecm,bbllx=54pt,bblly=380pt,bburx=558pt,bbury=727pt}} \\caption{LEFT: The distribution of the solar cycle length does not approach a Dirac function, the `quality' of the cycle only gives values of about 5. RIGHT: The wavelet spectrum of the sunspot-number time series shows two  peaks for both 10~yr and 100~yr (Frick {\\it et al.}, 1997a).} \\label{quality} \\end{figure} ",
        "conclusions": ""
    },
    "0201/astro-ph0201290_arXiv.txt": {
        "abstract": "Chandra deep fields represent the deepest look at the X--ray sky. We analyzed the Chandra Deep Field South (CDFS) with the aid of a dedicated wavelet-based algorithm. Here we present a detailed description of the procedures used to analyze this field, tested and verified by means of extensive simulations. We show that we can safely reconstruct the Log N--Log S source distribution of the CDFS down to limiting fluxes of $2.4\\times 10^{-17}$ and $2.1\\times 10^{-16}$ erg s$^{-1}$ cm$^{-2}$ in the soft (0.5--2 keV) and hard (2--10 keV) bands, respectively, fainter by a factor $\\gsim 2$ than current estimates. At these levels we can account for $\\gsim90\\%$ of the 1--2 keV and 2--10 keV X--ray background. ",
        "introduction": "The Chandra observatory is providing the astronomical community with the deepest X--ray look at the sky (Mushotzky et al. 2000; Hornschemeier et al. 2000; Giacconi et al. 2001). Two 1 Ms observations have been recently carried out: one in the northern hemisphere on the Hubble Deep Field North (HDFN) and the other in the southern hemisphere on a field, named Chandra Deep Field South (CDFS), selected for its low column density (the southern twin of the Lockman hole) and for the lack of bright X--ray and optical sources. The data are available from the Chandra public archive\\\\ ({\\tt http://asc.harvard.edu/udocs/ao2-cdf-download.html}) A further 1 Ms data set on the HDFN will be available in the next future.  The main goal of these observations is to look at the X--ray sky at the deepest level and to gain insight in the population of emitting sources comprising the cosmic X--ray background (XRB). The details of the data reduction and analysis procedures that have been applied to manage the CDFS data-set are discussed in Section 2. In order to fully exploit the potential of these data, refined detection algorithm have to be used. We developed a wavelet-based source detection algorithm (Lazzati et al. 1999; for its main characteristics see also Section 3), that we applied to the full sample of ROSAT HRI fields (Campana et al. 1999; Panzera et al. 2002, in preparation). We modified this detection algorithm, called Brera Multi-scale Wavelet (BMW), to account for the specific characteristics of the Chandra ACIS Imaging and Spectroscopic instruments. The algorithm (BMW-Chandra) has been extensively tested in the extreme conditions provided by the CDFS (Section 4). Our final goal is to obtain a source detection and a source Log N--Log S at the faintest limits in order to resolve as much as possible of the XRB in point sources. In order to compare our results with previous investigations, we carry out the analysis in the soft (0.5--2 keV) band and in the hard (2--10 keV) band. We are able to safely reconstruct the Log N--Log S source flux distribution down to $2.4\\times 10^{-17}$ and $2.1\\times 10^{-16}$ erg s$^{-1}$  cm$^{-2}$ in the soft and hard bands, respectively (Section 5). A first account of these results have been given in Campana et al. (2001), here we extend these results discussing in more details the resolved background and source characteristics. Conclusions are reported in Section 6. ",
        "conclusions": "Based on the experience developed in the ROSAT HRI data analysis (Lazzati et al. 1999; Campana et al. 1999) we build up a procedure for the automatic detection and characterization of sources in Chandra ACIS fields. The main difference between our wavelet algorithm (BMW-Chandra) and the standard CIAO {\\tt wavdetect} (Freeman et al. 2001) is the use of the multi-resolution technique. As explained, this makes the BMW-Chandra algorithm less flexible but faster and therefore more suitable for the analysis of large data sets and for Monte Carlo tests. The Chandra deep fields are, of course, the most intriguing fields to probe the algorithm. Here we have described the procedures we used to simulate and analyze in details the CDFS. The simulations we performed allow us to safely conclude that our detection procedures are robust in terms of number of spurious sources, positional accuracy, flux estimate down to the quoted limits. In this respect the Eddington bias correction is especially important, allowing us to approach the theoretical detection limit of 3 photons (Damiani et al. 1997b).  Scientific results have been previously discussed in Campana et al. (2001), here we improve and complement them. We independently detect 244 and 177 sources in the soft and hard band, respectively for a total number of 278 sources. Source distributions can be safely recovered down to fluxes $2.4\\times 10^{-17}$ and $2.1\\times 10^{-16}$ erg s$^{-1}$ cm$^{-2}$ in the soft and hard energy band, respectively. These are a factor of $\\gsim 2$ deeper than current estimates. Note that being still in the photon detection limit, an improvement by a factor of $\\sim 2$ in the Log N--Log S reconstruction corresponds to an effective doubling of the observing time.  There is general consensus that the main constituents of the soft 1--2 keV XRB can be ascribed to broad line AGN (i.e. Seyfert 1 galaxies; Hasinger et al. 1998; Schmidt et al. 1998). At faint fluxes ($\\lsim 10^{-15}$ erg s$^{-1}$ cm$^{-2}$) nearby ($z\\lsim 0.6$) optically-normal (possibly low luminosity AGN) galaxies are also being detected as soft sources (Fiore et al. 2000; Barger et al. 2001; Tozzi et al. 2001; Brandt et al. 2001a; Schreier et al. 2001; Koekemoer et al. 2001). Our analysis of the CDFS extends the source flux distribution down the faintest level of $2.4\\times 10^{-17}$ erg s$^{-1}$ cm$^{-2}$. We found that even at these low fluxes the Log N--Log S distribution can still be represented by the extrapolation from higher fluxes without upward trends. Including the contribution from cluster of galaxies (Rosati et al. 1998), we are able to resolve in point sources $>90\\%$ of the soft XRB, the exact value depending on the sky level itself. Our results are also in agreement with the recent fluctuation analysis on the HDFN down to $7\\times 10^{-18}$ erg s$^{-1}$ cm$^{-2}$ (Miyaji \\& Griffiths 2002).  In the hard band the XRB is made by the superposition of absorbed and unabsorbed sources (Setti \\& Woltjer 1989). The steeper power law index observed in bright AGN implies that a population of faint hard sources is present, likely affected by a considerable absorption. Studies with Chandra and XMM-Newton are now discovering these sources identifiable with Type II AGN unrelated to the morphological type (Schreier et al. 2001; Koekemoer et al. 2001). We note that only a small fraction of hard source are not detected in the soft band, likely indicating that at least $1-10\\%$ of the flux is in any case emitted in the soft band. Our analysis lead to the extension of the Log N--Log S source flux distribution down to $2.1\\times 10^{-16}$ erg s$^{-1}$ cm$^{-2}$. At this level we can rule out the UHURU/HEAO-1 estimate and account for $87\\%$ of the BeppoSAX/ASCA estimate. The fluctuation analysis by Miyaji \\& Griffiths (2002) does not add much stopping at $1\\times 10^{-16}$ erg s$^{-1}$ cm$^{-2}$ but with a large error.  The analyzed field of view is small and cosmic variance likely plays a role. We plan to carry out the same analysis of the Chandra HDFN in order to confirm the present results and to perform a joint analysis with other shallower and more extended surveys in order to cover with good statistics the bright flux ends (Moretti et al., in preparation).  Chandra with its superb angular resolution opens the possibility to extend at even lower fluxes the source distribution. In fact, at the present level sources are still photon limited. This will allow to deepen our knowledge of the distant universe, allowing to probe the accretion power over the history of the X--ray universe, and its implications for structure formation and the epoch of reionization."
    },
    "0201/astro-ph0201259_arXiv.txt": {
        "abstract": "I present photometric and radial velocity data for Galactic Cepheids, most of them being in the southern hemisphere.  There are 1250 Geneva 7-color photometric measurements for 62 Cepheids, the average uncertainty per measurement is better than $0.01^m$.  A total of 832 velocity measurements have been obtained with the CORAVEL radial velocity spectrograph for 46 Cepheids.  The average accuracy of the radial velocity data is 0.38 \\kms. There are 33 stars with both photometry and radial velocity data. I discuss the possible binarity or period change that these new data reveal.  I also present reddenings for all Cepheids with photometry.  The data are available electronically. ",
        "introduction": "A significant amount of photometry and radial velocity data for Cepheids had been published by \\citet{ber94a} and \\citet{ber94b}. These have been used to devise a new version of the Baade-Wesselink method \\citep{bbk97}.  However most Cepheids in Bersier, Burki \\& Kurucz's sample are in the northern hemisphere. New observations have been obtained for (mostly) long-period Cepheids visible from the southern hemisphere.  I present here these new observations, along with older unpublished data. The photometry is in the Geneva 7-color system and the radial velocity data have been obtained with the CORAVEL radial velocity scanner \\citep{bmp79}. ",
        "conclusions": ""
    },
    "0201/astro-ph0201545_arXiv.txt": {
        "abstract": "III~Zw2 was observed with XMM-Newton in July 2000. Its X-ray spectrum can be described by a powerlaw of photon index $\\Gamma \\approx 1.7$ with a Gaussian line at 6.7~keV. There is no significant evidence of intrinsic absorption within the source or of a soft X-ray excess. Multi-wavelength light curves over a period of 25 years show related variations from the radio to the X-rays. We interpret the radio to optical emission as synchrotron radiation, self-absorbed in the radio/millimeter region, and the X-rays as mainly due to Compton up-scattering of low-energy photons by the population of high-energy electrons that give rise to the synchrotron radiation. ",
        "introduction": "III~Zw2 (PG 0007+106, Mrk 1501) is a Seyfert I galaxy (Arp 1968; Khachikian \\& Weedman 1974; Osterbrock 1977) with z = 0.089 (de Robertis 1985). Superluminal motion of radio emitting plasma has been observed recently in the source, and this is the first detection of its kind in a spiral galaxy (Brunthaler et al.~2000). The source has long been known to show large-amplitude flux variations in the radio (Wright et al.~1977; Schnopper et al.~1978; Landau et al.~1980; Aller et al.~1985; Falcke et al.~1999) and the optical (Lloyd 1984; Clements et al.~1995). Variations of smaller amplitude (of less than 50~$\\%$), have also been detected in the IR (Lebofsky \\& Rieke 1980; Sembay, Hanson \\& Coe 1987) and the UV (Chapman, Geller \\& Huchra 1985). The X-ray temporal behaviour is less well studied, but comparison of observations at different epochs hints that the X-ray flux might vary substantially. The X-ray spectra of III~Zw2, obtained by SAS-3 (Schnopper et al.~1978), Ariel VI (Hall et al.~1981) and Einstein SSS (Petre et al.~1984), can be described by a powerlaw of photon indices $\\Gamma \\sim 1.3 - 1.7$, with neutral absorption consistent with that of the Galactic column. ",
        "conclusions": ""
    },
    "0201/astro-ph0201223_arXiv.txt": {
        "abstract": "\\nop We discuss how the space of possible cosmological parameters is constrained by the angular diameter distance function, $D_A(z)$, as measured using the SZ/X-ray method which combines Sunyaev-Zel'dovich (SZ) effect and X-ray brightness data for clusters of galaxies. New X-ray satellites, and ground-based interferometers dedicated to SZ observations, should soon lead to $D_A(z)$ measurements limited by systematic rather than random error. We analyze the systematic and random error budgets to make a realistic estimate of the accuracy achievable in the determination of $(\\Omega_m,\\Omega_\\Lambda,h)$, the density parameters of matter and cosmological constant, and the dimensionless Hubble constant, using $D_A(z)$ derived from the SZ/X-ray method, and the position of the first ``Doppler'' peak in the cosmic microwave background fluctuations. We briefly study the effect of systematic errors. We find that $\\Omega_m$, $\\Omega_\\Lambda$, and $w$ are affected, but $h$ is not by systematic errors which grow with redshift. With as few as 70 clusters, each providing a measurement of $D_A(z)$ with a $7\\%$ random and $5\\%$ systematic error, $\\Omega_m$ can be constrained to $\\pm 0.2$, $\\Omega_\\Lambda$ to $\\pm 0.2$, and $h$ to $\\pm 0.11$ (all at $3\\sigma$). We also estimate constraints for the alternative three-parameter set $(\\Omega_m,w,h)$, where $w$ is the equation of state parameter. The measurement of $D_A(z)$ provides constraints complementary to those from the number density of clusters in redshift space. A sample of 70 clusters ($D_A$ measured with the same accuracy as before) combined with cluster evolution results (or a known matter density), can constrain $w$ within $\\pm 0.45$ (at 3$\\sigma$). Studies of X-ray and SZ properties of clusters of galaxies promise an independent and powerful test for cosmological parameters. ",
        "introduction": "\\label{S:INTRO}   What set of cosmological parameters characterizes our Universe?   According to the most popular cold dark matter (CDM) scenario, the Universe consists of baryonic matter and a substantial amount of ``dark'' matter. A variety of recent measurements have led to the conclusion that the matter density parameter $\\Omega_m \\approx 0.3$ \\citep{Turn00}, while CMB measurements strongly favor a flat space time \\citep{Bernet01,Lee_et01}, and SNe Ia measurements indicate that the Universe is accelerating, suggesting a negative pressure \\citep{Rieset00,Perlet99}. Taken together, these pieces of evidence suggest that the baryonic and dark matter content of the Universe is supplemented by an additional smooth component with negative pressure, $P_w$, modeled by the equation of state $\\rho_w \\, c^2 = -w \\, P_w$, where $\\rho_w$ is the density of this component, $w$ is a dimensionless state parameter of order unity (cf. Huterer \\& Turner 2000).   Each existing dataset constrains, with limited accuracy, some subset of the cosmological parameters. Different measurements and combinations of measurements, such as SNe Ia, Cosmic Microwave Background (CMB) fluctuations, IRAS infrared galaxy surveys, classical double radio galaxy properties, 1.2-Jy galaxy redshift surveys, gravitational lensing, cluster X-ray temperature function and cluster number counts, baryon and gas mass fraction, and the SZ effect have been used to constrain cosmological parameters \\citep{Jaffet00,Balbet00,TegmZald00, Gueret00,Efstet99,Laseet99,Perlet99,GewiSilk98,Line98,Whit98,Pen_97,Sasa96, HuteTurn00,MajuSubr00,Bridet99,Dieget01}.   In the future, the SNAP project (http://snap.lbl.gov) plans to use the SNe Ia method to determine the matter density and the cosmological constant at the few percent level. Even with the next-generation CMB satellites, MAP and Planck, degeneracies will remain among the cosmological parameters that can be estimated from the results (Efstathiou \\& Bond 1999; Zaldarriaga, Spergel \\& Seljak 1997). The importance of using a wide range of methods, therefore, is twofold. First, a simultaneous consideration of all data sets should allow the best joint estimation of the cosmological parameters. Second, the agreement of different techniques for measuring the cosmological parameters should provide a cross-check of our understanding of the underlying processes and a control against systematic errors. As we extend our analysis of the CMB to more complicated models (tensor fluctuations, finite neutrino masses, etc.) the number of cosmological parameters increases, and it becomes even more important that the widest possible range of datasets is used, and that strong controls against systematic errors are in place.   Many of the techniques of cosmological parameter estimation use clusters as tracer particles. As a result there is a large number of planned cluster surveys in the two most important non-optical observational indicators of clustering: the SZ effect and cluster X-ray emission. Sunyaev-Zel'dovich effect surveys with dedicated interferometers or receiver arrays will observe hundreds of clusters with $z > 0.5$ per year (Browne et al. 2000; Holder, Carlstrom \\& Mohr 2000; Bartlett 2000). The new X-ray missions (Chandra, XMM) will provide data on hundreds of clusters with high redshift through their deep and medium-deep surveys.   Cluster evolution, the redshift distribution of clusters from SZ and X-ray surveys, $N_{SZ}(z)$ and $N_X(z)$, and cluster number counts as a function of X-ray flux, $N_X(S)$, are important constraints on cosmological parameters. While methods based on the CMB power spectrum and SNe Ia are sensitive to the angular diameter distance, cluster evolution (and number counts) is sensitive to the growth function of matter density fluctuations. \\cite{Bart00} estimated the performance of ground-based, arcminute-resolution, SZ surveys and concluded that more clusters will be detected with  deep, small-area surveys than shallow, wide-area surveys. \\cite{Kneiet01} studied the performance of the Arcminute MicroKelvin Imager experiment and showed that a set of only about 20~clusters, with redshifts in the range $z =$ 0 - 0.8 is needed to measure $N_{SZ}(z)$ sufficiently well to distinguish between $\\Omega_m$ = 1 and $\\Omega_m$ = 0.3 cosmologies. \\cite{Carlet01} discuss a deep SZ ground based survey, and quantify constraints from $N_{SZ}(z)$ on $\\Omega_m$ and $\\Omega_\\Lambda$. $N_{SZ}(S)$ and $N_{SZ}(z)$ were estimated from the proposed shallower, but all-sky Planck survey by \\cite{Dieget01}, who concluded that about 300 clusters (with the necessary optical follow-up to measure redshifts) would suffice to distinguish between open $\\Omega_m$ = 0.3 and flat $\\Omega_m$ = 1 cosmologies at $3\\sigma$ confidence. Holder, Haiman \\& Mohr (2001) discussed the constraints on the parameter space defined by $(\\Omega_m, \\Omega_\\Lambda, \\sigma_8)$ (where $\\sigma_8$ is the normalization of the matter power spectrum) using cluster evolution. Holder et al. showed that constraints from cluster evolution and SNe Ia observations are highly complementary to each other. \\cite{Haimet00} discussed the constraints on the $(\\Omega_m, w, h)$ parameter space, assuming a spatially flat geometry ($\\Omega_\\Lambda = 1 - \\Omega_m$), that follow from an SZ effect survey and a large angle deep X-ray survey (the Cosmology Explorer; Ricker and Lamb). They found that $N_{SZ}(z)$ and $N_X(z)$, combined with constraints from CMB or SNe Ia experiments, significantly reduce the degeneracies between $\\Omega_m$, $w$, and $h$. \\cite{HuteTurn00} estimated the constraints on $\\Omega_m$ and $w$ for flat geometry that can be gained by combining results from SNAP, Planck and SZ and X-ray surveys.    As has been realized, the angular diameter distance-redshift relation, $D_A(z)$, is at the heart of many of these techniques, and is sensitive to some important combinations of cosmological parameters while being degenerate under others \\citep{Jaffet00,TegmZald00,Efstet99,Laseet99,Perlet99,Whit98}. Recently \\cite{Whit98} estimated constraints on the pairs of quantities $(\\Omega_m,\\Omega_\\Lambda)$ and $(\\Omega_m,w)$ (the latter in a flat Universe) from the $D_A(z)$ function based on current SNe Ia data combined with CMB first peak constraints. The analysis shows that the constraint on parameters based on $D_A(z)$ is nearly orthogonal to the constraint based on the position of the first peak in the CMB fluctuation spectrum. These two datasets are thus highly complementary, and form a particularly powerful pair of measurements (see also Tegmark et al. 1998).   The shape and normalization of the observed angular diameter distance function constrains several cosmological parameters (the standard formulae for distance in Friedmann-Robertson-Walker Universes are given in, for example, Peebles 1993). The distance - redshift function, $D_A(z)$, in CDM models depends on the matter density, cosmological constant and Hubble constant, and any other particle density which contributes to the curvature of space-time. The slope of the distance-redshift function at low redshift is a measure of the Hubble constant, while the shape of the function depends on the curvature and the different densities. In Figure~\\ref{F:FIG1} we show the fractional difference in $D_A(z)$ with fixed matter density and Hubble constant ($\\Omega_m = 0.3$, $h = 0.65$), but various values cosmological constants ($\\Omega_\\Lambda$ = 0.7, 0.6, 0.3, solid, dashed and dash dotted lines) relative to a model with zero cosmological constant ($\\Omega_\\Lambda = 0$). It can be seen from this figure, that $D_A(z)$ is most sensitive to the value of the cosmological constant at redshift about unity, and quite insensitive to that at small or high redshifts. In the redshift interval from $z=0.5$ to $z=1.8$ the angular diameter distance for flat ($\\Omega_m=0.3$, $\\Omega_\\Lambda=0.7$) model is more than 10$\\%$ different from a model with the same matter density but zero cosmological constant ($\\Omega_m=0.3$, $\\Omega_\\Lambda=0$).   We can expect high precision data from hundreds of clusters of galaxies in the near future. With the present instrument suite, the statistical errors on individual measurements will be small, and so the usefulness of the data will be limited by their systematic errors. In this paper we evaluate the error budget of distance determination based on the SZ effect and X-ray measurements (assuming that the X-ray output is dominated by thermal bremsstrahlung, as is appropriate for the hot clusters in which the SZ effect is strong), and provide a realistic estimate of errors achievable in the angular diameter distance. We estimate how well one will be able to constrain the two parameter sets $(\\Omega_m, \\Omega_\\Lambda, h)$, and $(\\Omega_m, w, h)$ (assuming a spatially flat geometry, $\\Omega_\\Lambda = 1-\\Omega_m$) if this $D_A(z)$ function is combined with the position of the first Doppler peak in the angular power spectrum of CMB fluctuations.   Our treatment is complementary to previous work of \\cite{Kneiet01}, \\cite{Carlet01}, \\cite{Dieget01}, \\cite{Holdet01}, \\cite{Haimet00}, and \\cite{HuteTurn00}, who used $N_{SZ}(z)$, $N_{SZ}(S)$, and $N_{X}(z)$ to constrain cosmological parameters, since we discuss the importance of $D_A(z)$. It also complements the work by \\cite{Whit98}, who used $D_A(z)$ determined from existing SNe Ia data to constrain $(\\Omega_m, \\Omega_\\Lambda)$, and $(\\Omega_m, w)$ (flat): the errors from the SZ/X-ray technique have significantly different characteristics, and we are concerned with the limitations that will be encountered with future survey data.   In the next section we briefly describe the well-known method of angular distance determination based on the SZ effect and thermal bremsstrahlung with an emphasis on how measurements are used, and how their error propagate to the angular diameter distance. In section~\\ref{s:errors} we give a detailed analysis of the error budget, in section~\\ref{s:Constraints} we discuss the constraints on cosmological parameters. Finally, section~\\ref{S:Conclusion} summarizes our conclusions. ",
        "conclusions": "\\label{S:Conclusion}   We have estimated the accuracy achievable in the determination of cosmological parameters using the SZ/X-ray method of distance determination. This method uses a sample of clusters to map the distance-redshift relation, which is a sensitive probe of cosmological parameters. The advantages of this well-known method are: (1) unlike other distance determination methods, it depends only on the geometry of the Universe, and its average densities; (2) it is a physical method, based on relatively simple gravitational virialization of clusters (as opposed to complicated physics and chemistry involved in galaxy formation and supernova explosions); and (3) a large number of clusters is available for observation, and thus systematic effects can be reduced by using many clusters, or selecting clusters appropriately to reduce systematics. The necessary data should be available within the next few years.   The SZ/X method, as other cosmological tests, can constrain well only some combination of cosmological parameters. We have shown that constraints on ($\\Omega_m$, $\\Omega_\\Lambda$, $h$) from the $D_A(z)$ function measured in $z = 0 - 1$ are nearly orthogonal to constraints from the first peak of the CMB fluctuations, and also to constraints from cluster evolution. Constraints on ($\\Omega_m$, $w$, $h$) from the $D_A(z)$ function are complementary to those from cluster evolution (compare our Figure 2 to Figures~8 and 9 from Haiman et al. 2000). In general, $D_A$ provides constraints on cosmological parameters similar to those from the SNe Ia method. Constraints from cluster evolution (the source counts $N_{SZ}(z)$ and $N_{X}(z)$) are similar to constraints from the full CMB fluctuation spectrum as will be measured by MAP and Planck.  This result suggests that cosmological tests using {\\it only} clusters of galaxies can be an important independent check on the values of cosmological parameters measured by other techniques. Cluster-based methods have systematic errors which are unrelated to those of other methods of measuring cosmological parameters (e.g., from CMB fluctuations, and SNe~Ia). Furthermore, when the set of cosmological parameters is further extended, with additional components of density or structure parameters, joint analyses using multiple cosmological tests will be essential to remove the parameter degeneracy exhibited by any one test.   We demonstrated the effect of systematic errors on the determination of cosmological parameters. If random errors can be kept at a few percent level, systematic errors of similar magnitude will dominate the error in the Hubble constant. $\\Omega_m$, $\\Omega_\\Lambda$, and $w$ are not affected by redshift-independent systematic errors. We also showed that a systematic error with a linear gradient in redshift will not affect Hubble constant determinations, but causes systematic shifts in the estimates of $\\Omega_m$, $\\Omega_\\Lambda$, and $w$.  We showed, that in the near future, even with only 35 high redshift (and 35 low redshift) clusters, the SZ/X-ray method, combined with the position of the first peak in the power spectrum of CMB fluctuations, will provide enough accuracy to exclude $\\Omega_\\Lambda = 0$ models with high confidence even with constant and redshift dependent systematic errors (see Figure~\\ref{F:FIG4}). Also, the SZ/X-ray method, combined with the $\\Omega_m h^2$ = constant constraint, or cluster abundance, will allow us to determine the equation of state parameter, $w$, to within 0.45 (3$\\sigma$; Figure~\\ref{F:FIG5}), even with usual (redshift-independent) systematic errors.   This idealized discussion of the SZ/X-ray method and its errors in the determination of cosmological parameters could be improved by using the detailed characteristics of specific instruments and observing strategies. Departures from the assumptions made here could cause increases or decreases in the level of error in the derived parameters, with the largest changes likely for different redshift samplings.   We conclude, that the determination of the angular diameter distance - redshift function using the SZ effect and X-ray thermal bremsstrahlung emission from clusters of galaxies can be used with confidence to constrain cosmological parameters. In general, clusters of galaxies {\\it alone} can be used to constrain cosmological parameters independently from other methods. Clusters lead to parameter limits which are competitive with using other techniques.    Most of this work was done while SMM held a National Research Council Research Associateship at NASA Goddard Space Flight Center. We thank David Spergel and the anonymous referee for useful comments and suggestions."
    },
    "0201/astro-ph0201015_arXiv.txt": {
        "abstract": "The simple reading of the evidence is that the large elliptical galaxies existed at about the present star mass and comoving number density at redshift $z=2$. This is subject to the usual uncertainties of measurement and interpretation in astronomy, but should be taken seriously because it is indicated by quite a few lines of evidence. And it might be a guide to a more perfect theory of galaxy formation. ",
        "introduction": "Current ideas about structure formation suggest roughly half the large elliptical galaxies were assembled at redshifts less than unity. The measurements do not rule this out, as noted in many observational papers, but a far simpler interpretation is that the large ellipticals formed well before $z=1$.  The commonly discussed structure formation theory, $\\Lambda$CDM,\\footnote{This assumes the mass of the universe is dominated by cold dark matter, with gravitational growth of structure out of primeval adiabatic, Gaussian, and near scale-invariant density fluctuations in a low density cosmologically flat universe. I adopt Hubble parameter $H_o=70$ km~s$^{-1}$ Mpc$^{-1}$, matter density parameter $\\Omega =0.25$, zero space curvature, and a cosmological constant that makes fractional contribution $\\Omega _\\Lambda = 0.75$ to $H_o{}^2$.} certainly deserves careful attention because of its dramatic success in the interpretation of the temperature anisotropy of the 3K thermal cosmic background radiation. But the test is limited, and aspects of the theory seem problematic: small-scale structure (Moore et al. 1999 and references therein) and the void phenomenon (Peebles 2001). The case for early formation of the elliptical galaxies deserves careful attention too, because it is the simple interpretation of an impressively broad range of observations.  The main competing pictures of galaxy formation, at high redshift and low, often are termed monolithic and hierarchical. This makes historical sense,\\footnote{Thus the Partridge \\&\\ Peebles (1967) fit of the Friedmann-Lemaitre cosmology to the Eggen, Lynden-Bell \\&\\ Sandage (1962) picture for formation of the Milky Way indicates galaxies formed at $z\\sim 20$. Tinsley (1972) and Larson (1974) pioneered the modern picture for the properties of an elliptical that formed as a galaxy of stars at high redshift. At another extreme, Kauffmann, Charlot, \\&\\ White (1996) discuss evidence that only about one third of present-day early-type galaxies were assembled at $z=1$, roughly in line with what would be expected in $\\Lambda$CDM.} but can be confusing, because one can imagine galaxies form by hierarchical growth of structure either at high redshift or low. I will take ``formation'' to be the assembly of more than half the material now in the luminous parts of a galaxy, gathered within a sphere of radius 30~kpc, let us say. The epoch at which half the present-day stars formed may precede or follow assembly. Either way, galaxy evolution may continue well after formation as defined here. This includes merging of satellite galaxies, as in the Ostriker-Tremaine (1975) ``cannibalism'' picture, and more major mergers of galaxies that formed in tight pairs, as might be seen in clusters at high redshift (van Dokkum et al. 2000) and low (Struble \\&\\ Rood 1981). It also includes new generations of stars from gas from the outer parts of the galaxy, or recycled from evolved stars in the ``frosting'' picture of Trager et al. (2000).  I limit this discussion to the formation of large elliptical galaxies, or in some cases E plus S0 galaxies. There is a lot to be learned from late-type galaxies, of course, but the limitation helps focus the discussion on systems that present us with a fascinating pattern of regularities. The simplicity of the early-type galaxies allows close exploration of rich details of departures from the regularities. Both aspects, regularity and complexity, may prove to be important guides to how and when the galaxies formed and what that might teach us about the physics of our expanding universe.  Since I am more taken by the side of simplicity I should acknowledge here that this certainly is not the whole story: the departures from the patterns reflect significant differences among the histories of galaxies (Toomre 1977; Schweizer 2000). We see this in strongly disturbed late-type systems, whose remnants likely will qualify as new early-type galaxies, in more moderately disturbed ellipticals, that have gained mass by recent accretion or mergers, and in Butcher-Oemler (1984) galaxies, that show us morphological transformations. Aspects of this complexity are reviewed in \\S 2.2.  The issue for the present discussion is whether the examples of late formation show us the main way galaxies formed, or amount to a perturbative effect on physical properties that were established much earlier. Arguments of some years' standing for the latter are that one sees large elliptical galaxies of old stars at high redshifts (Oke 1971, 1984), and that subsequent formation might be expected to erase the correlations of color and heavy element abundance with luminosity (Faber 1973; Visvanathan \\&\\ Sandage 1977; Ostriker 1980). The known patterns of regularities among early-type galaxies are much richer now. I devote most of this contribution to a review of the literature of these patterns, because I find them fascinating and surely educational. Just what we might be learning is briefly considered in \\S 3. ",
        "conclusions": ""
    },
    "0201/astro-ph0201365_arXiv.txt": {
        "abstract": "{We report SIMBA 1.2\\,mm dust continuum observations of the environments of eight methanol maser sources, all discovered during spatially fully-sampled, untargeted surveys of the galactic plane. We summarise our search for possible associations of the masers with IR sources (IRAS and MSX) and find that it is not always possible to make definite associations. A preliminary characterisation of the IR sources found in the maser neighbourhood is given according their position in the [60-25] -- [25-12] colour-colour diagram. ",
        "introduction": "Massive star formation has attracted increasing interest in recent years, with the mounting evidence that methanol masers provide excellent tracers and probes of the high-mass star formation process (e.g. Hunter et al. 1998; Minier et al. 2001). Such evidence has led to several surveys aimed at identifying new high-mass stars, often targeting IRAS colour-selected UCHII regions (Caswell et al. 1995; Schutte et al. 1993; Szymczak et al. 2000). Although making significant contributions to the discovery of new massive star forming regions (MSFRs), such surveys have yielded only a $\\sim$15\\% detection rate of methanol masers.  An alternative approach for discovering MSFRs has been provided by {\\it blind} surveys of 6.7\\,GHz methanol masers. In particular, the surveys of Ellingsen et al. (1996) and Pestalozzi et al. (2002, in prep.) have sampled in a consistent manner large regions of the galactic plane from the southern and northern hemispheres respectively. Since the majority of methanol masers revealed in these surveys does not have a clear association (within the positional accuracy) with IRAS sources, the question has arisen: could they be tracing early stages of the high-mass star formation process, in which the still highly-embedded, nascent star is radiating much of its flux in the millimetre and submillmetre regions (as suggested by Walsh et al. 1999)?  With the aim of answering this question, in addition to ascertaining the SEDs for regions for which there is existing IR data (IRAS and MSX), we report here on 1.2\\,mm bolometer observations of the regions surrounding eight 6.7\\,GHz methanol masers detected in blind surveys. This constitutes the first part of our follow-up investigations for understanding the methanol maser phenomenon, and the evolutionary route to forming massive stars.  \\vskip -1.5cm ",
        "conclusions": "  \\begin{itemize}  \\item Blind surveys of methanol masers facilitate the detection of young, still deeply embedded high-mass star formation regions; \\item IRAS sources associated with methanol masers show a flat density profile ($0 \\le p \\le 0.5$), which means that they can be classified as objects in an early stage of massive star formation.  \\end{itemize}"
    },
    "0201/astro-ph0201479_arXiv.txt": {
        "abstract": "We present global 3D MHD simulations of geometrically thin but unstratified  accretion disks in which a near Keplerian disk rotates between two bounding regions with  initial rotation profiles that are stable to the MRI. The inner region models the boundary layer between the disk and an assumed more  slowly rotating central, non  magnetic star. We investigate the dynamical evolution of this system in response to initial  vertical and toroidal fields imposed in a variety of domains contained within  the  near Keplerian disk. Cases with both  non zero and zero net magnetic  flux are considered and  sustained dynamo activity found in  runs for up to fifty orbital periods at the outer boundary of the near Keplerian disk.  We find a progression of behavior regarding the turbulence resulting from the MRI and the evolving structure of the disk and boundary layer according to the initial field configuration. Simulations starting  from  fields with small radial scale and with zero net flux lead to the lowest levels of turbulence and smoothest variation of disk mean state variables. As  found in local simulations, the final  outcome is shown to be independent of the form of the imposed field. For our computational set up, average values of the Shakura \\& Sunyaev (1973) $\\alpha$ parameter in the Keplerian disk are  typically $0.004\\pm 0.002.$ Magnetic field eventually always diffuses into the boundary layer  resulting in the build up of toroidal field, inward angular momentum transport and the accretion of disk material. The mean radial velocity, while exhibiting large temporal fluctuations is always subsonic.  Simulations starting  with net toroidal  flux may yield an average $\\alpha \\sim 0.04.$ While being characterized by one order of magnitude larger average $\\alpha$, simulations starting from vertical fields with large radial scale and net flux may lead to the formation of persistent  non-homogeneous, non-axisymmetric magnetically dominated regions of very low density. In these gaps, angular momentum transport occurs through magnetic torques acting between regions on either side of the gap. Local turbulent transport occurs where the  magnetic field is not dominant. These simulations are indicative of the behavior of the disk when  threaded by magnetic flux originating from an external source. However, the  influence  of such presumed  sources  in determining the boundary conditions  that should be applied to the disk remains to be investigated. ",
        "introduction": "\\label{S0} \\noindent The study of the boundary layer, the region where the angular velocity of  an accretion disk drops to match the rotation velocity of the central object, is of great importance for the understanding of accreting objects, since up to half of the total accretion energy  can be released there (Lynden-Bell $\\&$ Pringle 1974)  if the central object is a slow rotator. Up to now detailed studies of this region (eg. Papaloizou \\& Stanley 1986, hereafter PS, Kley 1989, Popham \\& Narayan 1992)  have been based on the Navier-Stokes equations with a modified viscosity prescription involving an anomalous viscosity coefficient. This is presumed to contain the effects of any turbulence present reducing the problem to  one of laminar flow.  \\noindent The decrease of angular velocity in the boundary layer is associated with an increase in the thermal pressure, which replaces  centrifugal support for the accreting  matter. The large pressure gradients in turn, may be associated with supersonic radial infall velocities in this region, if a standard disk viscosity prescription is  continued into the layer. However, it has been argued that if such a supersonic flow occurred, the star would loose its causal connection to the outer disk (Pringle 1977). In relation to this issue, several studies based on the Navier-Stokes equations  were performed but with a modified viscosity prescription in the boundary layer corresponding to its much lower pressure scale height (PS, Popham  \\& Narayan 1992). This modification alone  does not eliminate supersonic flows under all conditions (e.g. for large values of $\\alpha\\simeq 1$). Popham  \\& Narayan (1992)  suggested reducing the viscosity coefficient to zero as the radial infall velocity approaches the  sound speed. This causally limited viscosity  (Narayan 1992) always leads to subsonic infall. A different approach,  appreciating the short time required for matter to pass through the layer, allows the viscous stress components to relax towards their equilibrium values on a relaxation time scale (Kley \\& Papaloizou 1997) and so naturally incorporates causality. Studies of  one-dimensional models led to the conclusion that the boundary layer must be characterized not only by the value of $\\alpha$ in the outer disk, but also by the nature of the viscous relaxation process. Additional time dependent studies demonstrated that for low $\\alpha\\simeq 0.01$, the boundary layer adjusts to a steady state, while for large $\\alpha=0.1$, significant disturbances occurred in the boundary layer and to the power output. Periodic oscillations were seen throughout the disk by PS, whenever $\\alpha$ was large (close to 1). For small $\\alpha$, the oscillations were localized near the outer disk boundary. These oscillations are caused by the viscous overstability found by Kato (1978) and Blumenthal et al. (1984) and it was suggested they may be important in explaining the time-dependent behavior in accreting objects such as CVs and protostars.  \\noindent All of the above models were based on an ad-hoc anomalous viscosity  prescription  and  did not consider further  its origin. The discovery of the relevance of the magnetorotational instability (MRI) (Balbus  \\& Hawley 1991) has opened up a new era in  accretion disk astrophysics. The instability provides a robust and self-consistent mechanism for the production of turbulence and angular momentum transport in these objects if they are adequately ionized, thereby removing the need for  ad-hoc prescriptions. The development of different numerical codes has enabled a detailed investigation of the nonlinear phase of the instability. First numerical studies were performed in a local shearing box approximation, (Hawley, Gammie \\& Balbus  1995, Brandenburg et al. 1995, Sano, Inutsuka \\& Miyama 1998, Fleming, Stone \\& Hawley 2000, Miller  \\& Stone 2000). These studies show that the turbulent outcome of the MRI depends on the initial field configuration applied to the disk. Thus, local simulations with initial vertical fields with zero net flux field indicate an average  $\\alpha=0.001-0.006$, while vertical fields with non zero net flux result in larger values of $\\alpha$ up to $0.3.$ The turbulent outcome of an initially  unstable toroidal field can  lead to intermediate values of $\\alpha$ up to $ 0.04$ depending on the net flux.  \\noindent Relatively recently  studies of  instabilities in global  disk models have begun (Armitage 1998). Such studies are needed to see how an unstable disk   modifies its  underlying  structure  in response to globally varying levels of turbulence and whether the longer time scale evolution  is in any way like that of standard $'\\alpha'$ disk  models.  Recent  studies have  been made by Hawley (2000), who concentrated on the evolution of thick tori, and Hawley \\& Krolik (2001). The latter study focuses on the evolution of the inner region of a disk accreting onto a black hole modeled with a pseudo-Newtonian potential.  \\noindent In this paper we study the interaction of an accretion disk with a boundary layer region located between it and the central star. This situation is the relevant one to consider for non relativistic accretion onto a non magnetic central star. The inner boundary layer region together with an exterior stable  region also provide convenient, relatively inert regions in which to embed a near Keplerian disk with an unstable rotation profile. Instead of prescribing the viscosity in an ad-hoc fashion, as in previous studies of the boundary layer, we self-consistently incorporate the turbulence arising from the MRI as the source of viscosity and diffusion of magnetic field.  \\noindent We assume the disk to have a small ratio of  scale height to radius ( $H/R\\simeq 0.1$ ). The gravity is assumed to be entirely due to the central object and for simplicity and in common with other global disk studies we neglect vertical stratification  by  adopting  a cylindrically symmetric potential thus focusing the study on the radial structure of the disk. We study  the dynamical evolution  of the disk over a time span of up to one thousand rotation periods  measured at the inner disk  edge. We consider different initial magnetic field configurations (poloidal and toroidal) imposed in the main body of the disk. Cases with both  non zero and zero net magnetic flux are considered.  \\noindent Simulations starting with small scale initial fields with zero net flux exhibit the lowest Shakura  \\& Sunyaev (1973) parameter $\\alpha$ with  a   mean value averaged over the Keplerian domain  of $\\simeq 0.005$. In this case the simulations  on average attain a final state characterized by the same mean $\\alpha$ and  magnetic energy  independently of (within computationally defined limits) the initial field strength.   There may also be a tendency for  the mean value of $\\alpha$ to increase with the extent of the vertical domain and the numerical resolution. This  is in agreement with results of shearing box simulations (Hawley, Gammie \\& Balbus 1996, hereafter HGB96). On the other hand simulations with large net  magnetic flux may evolve turbulence with a larger mean value of $\\alpha$ of $\\simeq 0.04$, when the initial field is toroidal.  Models starting from  initial vertical  fields with large radial scale are such that $\\alpha$ attains maximum values $>1$ in regions associated with prominent density depressions, while outside these gaps, $\\alpha$ reaches values similar to the zero net flux models (as low as $0.005).$ The volume averaged $\\alpha$ depends on the initial plasma beta in such cases.  \\noindent We find that  all simulations locally  exhibit strong variations of the vertically and azimuthally averaged values of  $\\alpha$ in time and with  radius. All of the models display oscillations of the vertically and azimuthally averaged  radial Mach number. Even though the boundary layer region is stable to the MRI, magnetic field always diffuses into it. This is the case even when the initial field is non zero only well away from the layer. Toroidal field build up enables mass to accrete through it onto the central star through the operation of magnetic torques.  \\noindent The paper is organized as follows: in \\S \\ref{S1} we present the basic model and computational set up.  In \\S \\ref{S2} we describe the numerical procedure. In \\S  \\ref{S3} we discuss azimuthal and vertical averaging together with  the global transport of angular momentum and energy dissipation in the disk. \\S \\ref{S4} is devoted to the investigation of  cases with initial fields with zero net magnetic flux. In \\S \\ref{S5} we present  results when the initial magnetic field  has net flux. Finally, we summarize our results in \\S \\ref{S6} . ",
        "conclusions": "\\label{S6} \\noindent In this paper we have studied the  time dependent evolution of a near Keplerian accretion disk which rotates between two bounding regions with  initial rotation profiles that are stable to the MRI. The inner region models the boundary layer between the disk and a central star. Because we  self-consistently incorporate MRI-induced turbulence as the source of viscosity, the  necessity of an ad-hoc viscosity prescription of the type used in earlier studies  is removed.  \\noindent For this first study we assumed the disk to be unstratified with aspect ratio $ H/R \\simeq 0.1-0.2 $  by adopting a  cylindrically symmetric  potential  assumed to be exclusively generated by the central object. For most of the cases considered, the dynamical evolution of the disk following from the imposition of different magnetic field configurations was followed over a time span  between one hundred and two hundred rotation periods at the inner edge of the computational domain. However, simulations of radially more extended disks were followed up to one thousand orbital periods at the inner edge of the computational domain. For initial conditions both toroidal and poloidal magnetic fields with zero as well as net flux were applied in varying domains contained within the  near Keplerian disk.  \\noindent Simulations starting from toroidal and poloidal fields with zero net flux and a small scale of radial variation evolved to a  state characterized by a smooth angular velocity and density profile similar to the initial one. This was independent of the type, and within numerically determined limits,  the amplitude of the initial field. This result also holds in shearing box simulations (HGB96). But one must bear in mind that in numerical computations, the range of initial plasma betas to which one has reasonable access is restricted by the resolution. Typical values of $\\langle \\alpha \\rangle$  representing a volume average over the Keplerian  domain are 0.004$\\pm$0.002. Moreover, runs with a radially extended disk showed  that the saturated turbulent state is maintained over more than fifty orbital periods at the outer boundary of the Keplerian domain.  \\noindent While the shearing box approach guarantees flux conservation for all time in a Keplerian domain, in global simulations this is not necessarily the case. This is because in large scale simulations the inevitable diffusion of the magnetic field out of the Keplerian domain and into the boundary domains can lead to the violation  of flux conservation in the Keplerian domain. This happens even though the boundary domains are stable to the MRI. However, we found that different prescriptions of the inner  boundary layer region  do not affect the final state in the Keplerian domain even though the behavior of the boundary layer itself may   vary significantly. Once significant field has leaked into the boundary layer, toroidal field is built up due to the shear and causes $\\alpha$ to attain negative values corresponding to inward angular momentum transport and mass accretion. In such cases, the boundary layer expands outwards but radial motions remain subsonic.  \\noindent We also find a similar approximate correlation between the Keplerian domain volume averaged ($R \\phi$)-component of the stress and magnetic energy as  found in  local simulations:  For simulations with zero net flux we find, after averaging out short term variations, $\\langle \\alpha \\rangle\\sim  0.5/\\langle \\beta \\rangle$.  \\noindent Models with an initial vertical field with zero-net flux and large scale of radial variation exhibit local minima in the density associated with maxima of the angular velocity. These density pockets can reach a contrast of one order of magnitude with respect to the surroundings. The turbulent state can nonetheless be characterized by an average $\\alpha$ similar to the cases with small scale of radial variation. It may require a  very  long time scale  for such states to relax to those found in the initially small scale field cases.  \\noindent All models display large variations of $\\alpha$ in time and radius including oscillations on a rotational time scale. Variations of  one order of magnitude are typical.   \\noindent Models staring with fields that have non zero net flux  lead to a higher level of turbulence. Thus, model  n4  that started with a toroidal field attains an average alpha of 0.04. Those starting with vertical fields with net flux such as n1 may display several gaps in density, with the radially innermost gap typically located next to the boundary layer. The density contrast in the most prominent gaps can reach up to 2 orders of magnitude (in an azimuthal and vertical average) with respect to their surroundings. In the gap-regions, $\\alpha$ alternates between values $<1$ and $>1$, sometimes exceeding 3 and dropping to an average of 0.005 (comparable with the zero net flux simulations) in non gap regions. Values of $\\alpha$ exceeding 1 indicate angular momentum transport by magnetic torques originating from fields that connect across the gap region.  \\noindent Recognizing that there are issues to be resolved regarding the correct boundary conditions required to represent the effects of external conducting material, these solutions might be relevant when the disk becomes dominated by an external magnetic field. Such magnetic fields may affect a variety of processes that take place in accretion disks, from dust coagulation to the interaction of planets with the disks in which they have been formed.  \\noindent Clearly there is much room for future improvements and developments. Convergence needs to be checked at much higher resolution than currently attainable. Studies of  more extended inner MRI-stable boundary layers should be carried out, and vertical stratification should be included. In this connection the simple periodic boundary conditions used in the vertical direction leave the vertical flux relatively unconstrained and unconnected to any conductors external to the disk. Proper matching of boundary conditions to external fields is also an issue that may affect the behavior of the low density gap regions studied here and a subject for future study.     \\subsection{Acknowledgements}  We would like to thank Richard Nelson for encouragement and support regarding computational matters and him together with  Caroline Terquem  and Greg Laughlin for valuable discussions. We  acknowledge support from the UK Astrophysical Fluids Facility and the NASA Advanced Supercomputing Facility's Information Power Grid Project's Pool at NASA ames Research Center. AS thanks the Astronomy Unit at QMW for hospitality, the European Commission for support under contract number ERBFMRX-CT98-0195 (TMR network ``Accretion onto black holes, compact stars and protostars'') and the NRC  for a research  fellowship."
    },
    "0201/astro-ph0201153_arXiv.txt": {
        "abstract": "In this article we consider the growth of seed black holes immersed in dark matter halos. We first investigate the adiabatic growth in various initial distribution functions (isothermal, power law, and NFW) and find the resulting density, radial velocity, and anisotropy profiles. In addition we estimate the growth rate for a given black hole mass in the corresponding adiabatically modified dark matter distribution function. Only in the isothermal case is there a convincing black hole mass-age relation. By calculating the line of sight velocity dispersion for the various cases as a function of the black hole mass, we find the predicted adiabatic $M_{bh}-\\sigma$ relation; this never approaches the recently observed power law. We conclude by abandoning adiabaticity, suggesting that the black hole grows proportionally to the dark matter halo itself on a dynamic time scale. This allows us to relate the observed $M_{bh}-\\sigma$ relation to the cosmological power spectrum on galactic scales by using dimensional scaling arguments. ",
        "introduction": "Recent results by \\citet{geb00} and by \\citet{fer00} that establish a strong correlation between central black hole mass $M_{bh}$ and the velocity dispersion $\\sigma_e$, measured at $r_e/8$, indicate that the central black hole is intimately related to the dynamical structure of the galaxy. An earlier result by \\citet{mag98} relates $M_{bh}$ linearly to the mass of the bulge, which suggests a similar conclusion, while Merrifield, Forbes, \\& Terlevich (2000) establish a link between $M_{bh}$ and the age of the stellar system: a massive central black hole seems to grow over a period of a Gyr or so.  These results have already stimulated the emission of various theories that create a feed-back mechanism between bulge star formation and accretion of gas onto the black hole \\citep{bur01}. These theories generally tend to establish the desired correlations but at the cost of some rather complicated physics simply, perhaps oversimply, described.  The object of the present paper is to first revisit the correlations established by the adiabatic growth of a black hole in a galaxy, since this process is relatively free of physical assumptions once the initial distribution function (DF) is chosen. This part of the work is very much in the spirit of \\citet{mar99}, who examined the effect of the black hole on the fundamental plane relations, but did not consider explicitly the $M_{bh} - \\sigma$ relation. We choose distribution functions moreover that are appropriate for collisionless dark matter halos, the philosophy being that these are the dynamically dominant components of massive galaxies and should therefore dictate the observed velocities. These are chosen to be an isothermal or Gaussian DF (to test our code against previous work and because this may be the maximum entropy state according to \\citet{nak00}), an isotropic steady-state power law DF found by \\citet{hen95} for comparison purposes, and the DF that corresponds to the NFW density profile, as approximated by \\citet{wid00}, as the best measured approximation to a dark-matter halo.  Such an approach does not, of course, explain the origin of the black hole, but rather yields only the perturbed DF that is created by its adiabatically established presence. An earlier attempt to grow the black hole during the formation of the galaxy \\citep{ost00} relied on dissipative dark matter which is fraught unfortunately with badly known parameters. In the present approach the black hole growth is limited to the particle flux across the event horizon of a seed black hole that is peculiar to the initial or adiabatically modified DF.  The adiabatic approach fails to yield either the correct form or extent of the observed correlations between the black hole mass and the modified galaxy. We therefore suggest a possible explanation of the correlations based on a non-adiabatic process of black hole growth on the formation timescale of the dark matter halo. The argument is essentially dimensional at this stage and must be examined numerically in greater detail.  In section 2 we review the adiabatic growth approximation and verify our code with the isothermal DF. In section 3 we give the results in the power law and NFW distribution functions. In section 4 we discuss the resulting $M_{bh}$ versus $\\sigma_e$ in the various distribution functions and show the extent of the black hole influence in the galaxy. In section 5 we give our dimensional derivation of the observed relations based on a self-similar (but not necessarily spherically symmetric) growth of black hole and dark halo. Finally we give our conclusions. ",
        "conclusions": "In this paper we have explored the implications of growing a black hole in various dark matter halo distribution functions. Our principal approach was to assume that the black hole grows adiabatically on a time scale long compared with the dynamical time of the halo. The method gives definite predictions for the modified density, radial velocity and anisotropy profiles in the isothermal, self-similar ``power law'' and NFW dark matter distributions. The isothermal calculations reproduced and extended slightly previous work, but the calculations for the self-similar and NFW dark matter halos are new.  Depending on the mass of the black hole the disturbances can be noticeable out to nearly the core radius of the galaxy. Moreover, estimates of black hole growth time scales in the adiabatically modified distribution functions are given. Only in the isothermal DF is there found a reasonable black hole mass-galactic age relation. In no case however does the adiabatic argument give an $M_{bh}-\\sigma$ relation that  is close to that observed.  Thus in the concluding section we explored by dimensional argument based on the concept of multidimensional self-similarity \\citep{car91,hen97} the possibility that the central black hole grew on a dynamical time scale with the dark matter halo. A simple argument predicts $a=3/(1-\\epsilon/2)$ where $a$ is the observed power in the $M_{bh}-\\sigma$ relation and $-\\epsilon$ is the power in $r$ of the initial cosmolgical density perturbation that produced the galactic halo. Under certain assumptions this power can in turn be related to the power $n$ of the primordial cosmological power spectrum on the scale of galactic halos, so that reasonably, $a=12/(1-n)$. This gives $n=-2$ for $a=4$, which is close to that observed in both cases. We conclude that this suggestion is promising and more work should be done to confirm or infirm the idea that the black hole can grow self-similarly with the dark matter core."
    },
    "0201/astro-ph0201275_arXiv.txt": {
        "abstract": "We calculate the column density distribution of damped \\lya systems, modeled as spherical isothermal gaseous halos ionized by the external cosmic background. The effects of self-shielding introduce a hump in this distribution, at a column density $\\nhi \\sim 1.6\\times 10^{17} X^{-1} \\cm^{-2}$, where $X$ is the neutral fraction at the radius where self-shielding starts being important. The most recent compilation of the column density distribution, by Storrie-Lombardi \\& Wolfe, shows marginal evidence for the detection of this feature due to self-shielding, suggesting a value $X\\simeq 10^{-3}$. Assuming a photoionization rate $\\Gamma\\simeq 10^{-12} \\seg^{-1}$ from the external ionizing background, the radius where self-shielding occurs is inferred to be $\\sim 3.8 \\kpc$. If damped \\lya systems consist of a clumpy medium, this should be interpreted as the typical size of the gas clumps in the region where they become self-shielding. Clumps of this size with typical column densities $N_H\\sim 3\\times 10^{20} \\cm^{-2}$ would be gravitationally confined at the characteristic photoionization temperature $\\sim 10^4$ K if they do not contain dark matter. Since this size is similar to the overall radius of damped \\lya systems in Cold Dark Matter models, where all halos are assumed to contain similar gas clouds producing damped absorbers, this suggests that the gas in damped absorbers is in fact not highly clumped. ",
        "introduction": "Damped \\lya systems (hereafter, DLAs) contain most of the atomic hydrogen in the universe (e.g., \\citealt{Wolfe95,Storrie00}). They are therefore fundamental objects to understand galaxy formation: in the process of forming galaxies, the matter that is originally in the low-density, ionized intergalactic medium must invariably become first atomic, before it can form molecular clouds and stars.  The nature and geometrical shape of the DLAs is still a subject of debate. The two alternative hypotheses that are being discussed at present and confronted with the observations are as follows: (a) They are rotating disks, which may or may not be forming stars; (b) They are approximately spherical halos, where global rotation contributes little to the support against gravity. The truth may be a combination of these two possibilities: DLAs might be described as flattened halos or thick disks, with a range of the ratio of rotation velocity to velocity dispersion.  A fundamental property of DLAs is that their associated metal lines reveal the presence of multiple absorbers in most of the systems, which are thought to arise from clumps of gas. The multiple absorption lines appear over velocity intervals of $30$ to $300 \\kms$. The statistical properties of the multiple absorption lines have been used to attempt to distinguish between the disk and halo models, but both alternatives seem to be consistent with observations (\\citealt{Prochaska97,Prochaska98,Haehnelt98,Haehnelt00,McDonald99}). Although the observations reveal that there are usually one or a few clumps intersected along a random line of sight, with characteristic column densities of $\\sim 10^{20} \\cm^{-2}$, the size of the clumps is not well known. Constraints in a few systems from double lines of sight in lensed QSO's, and from model calculations of the ionization parameter, indicate sizes or structure in the range of 20 pc to 1 kpc (\\citealt{Rauch99,Lopez99}). Models of structure formation predict that the overall size of DLAs is of the order of 1 to 10 kpc (e.g., \\citealt{Katz96,Gardner97,McDonald99}). If the clump size is not much smaller than the size of DLAs, then the absorption lines would actually be arising from mild fluctuations in density and velocity in a turbulent medium; whereas if the clump sizes are much smaller, they would be real separate entities with a large overdensity relative to an interclump medium.  This paper discusses the effects of self-shielding on the form of the column density distribution of damped \\lya systems, and the dependence of this distribution on the size of the gas clumps. As the gas becomes self-shielding against the external cosmic ionizing background, there is a rapid transition from ionized to fully atomic gas, and this causes a feature in the column density distribution. If $X$ is the neutral fraction of the gas seen at a column density $\\nhi = 1.6\\times 10^{17} \\cm^{-2}$ (where the optical depth to photons at the hydrogen ionization edge is equal to one), a hump in the column density distribution will occur at $\\nhi\\sim 1.6\\times 10^{17} X^{-1} \\cm^{-2}$. Therefore, a measurement of $X$ from this feature in the column density distribution, combined with an independent estimate of the ionizing background intensity, can be used to infer the size of the gas clumps.  Observationally, compared with the expected number extropolated from low column density distributions, an excess of absorption systems with $\\nhi \\gtrsim 2\\times 10^{20} \\cm^{-2}$ is found by \\citet{Lanzetta91}. The recent compilation of the column density distribution of known DLAs by \\citet{Storrie00} shows more clearly a hump around $\\nhi\\sim 2\\times 10^{20}\\cm^{-2}$. Theoretically, based on an approximate calculation of the self-shielding effect, \\citet{Murakami90} showed that a flat part appears in the column density distribution. With a simplified consideration on the ionizing flux transfer, \\citet{Petitjean92} investigated models of self-gravitating, photoionized, spherical gaseous cloud and explained the flattening of the column density distribution at $\\nhi \\sim 2\\times 10^{20} \\cm^{-2}$ as the effect of self-shielding. Similar calculations of the expected column density distribution due to self-shielding have been done by \\citet{Corbelli01} for plane-parallel geometry.  Here, we will consider a spherical geometry with a singular isothermal profile to perform a self-shielding calculation, and focus on the application of inferring the size of the gas clumps in damped \\lya systems. ",
        "conclusions": "We have shown that self-shielding should cause a sharp transition from an ionized to an atomic medium in any gas cloud photoionized by the external cosmic background. The expected absence of photons above frequency $4\\nu_L$ due to \\heii absorption makes that transition even sharper once the neutral fraction has been increased by a factor 64 by self-shielding. This produces a hump in the column density distribution. The column density at which the hump occurs measures the quantity $\\Gamma r_{ss}$.  Most damped \\lya systems are likely to be photoionized by the external cosmic background. In fact, since their rate of occurrence is about one fourth per unit redshift at $z\\sim 3$ \\citep{Storrie96}, and the mean free path of ionizing photons is $\\Delta z \\simeq 1$ \\citep{Miralda90}, it is easily shown that if all the sources of the cosmic ionizing background were embedded inside damped \\lya systems, then the local source in any such system would contribute about the same flux as the external background on average. It is much more likely that any sources associated with typical damped \\lya systems do not contribute significantly to the cosmic background.  The results for the column density distribution presented by \\citet{Storrie00} suggest a possible detection of the self-shielding effect, which would imply a radius $r_{ss} \\simeq 3.6 \\kpc$ for $\\Gamma = 1.25\\times 10^{-12} \\seg^{-1}$. This represents a measurement of the size of the clumps in damped \\lya systems. Even though our calculation is for a spherical, isolated cloud, a system of randomly located clumps would produce a similar column density distribution depending on the radius of the clumps, as long as each clump is illuminated by the ionizing background with an optical depth $\\lesssim 1$ along a large fraction of directions. This should be correct since the typical number of intersected clumps identified in the metal absorption lines is not very large.  The typical size of damped \\lya systems inferred from CDM models where all halos with $V_c \\geq 40 \\kms$ give rise to the absorption systems is 1 to 10 $\\kpc$ (e.g., \\citealt{Katz96,Gardner97,McDonald99}). Hydrodynamic cosmological simulations show that, at redshift $z\\sim 2-4$, for halos more massive than $\\sim 10^{11}\\msun$, the projected distance of DLAs to the center of the nearest galaxy is around $10 \\kpc$ \\citep{Gardner01}. Since the clump size we infer is about the same, this suggests that the multiple metal absorption lines do not arise from highly overdense clumps, but from mild density fluctuations in gaseous halos. If there were highly overdense clumps, their small self-shielding radii would produce a column density distribution with a less pronounced hump at lower column density than the curves in Figure 2. The appearance of individual absorption lines in the spectra of \\mgii and other metal lines might be deceiving, and could be due to a continuous medium, where the \\mgii density is highly sensitive to the gas density due to photoionization and self-shielding effects.  Nevertheless, a number of caveats must be borne in mind in interpreting the observed column density distribution as a measurement of $\\Gamma r_{ss}$. In reality, there should be a distribution of clump sizes, and differences from a simple spherical geometry would introduce additional variations, which would smooth the shape of the distribution predicted in Figure 2. Furthermore, the density profile of the gas cloud can be different from what we assume here. A fully three-dimensional calculation of self-shielding in a hydrodynamic simulation of galaxy formation would be highly desirable for a more robust interpretation of the observations of damped \\lya systems.  Finally, we mention that once the size of the clumps is known, their column densities can be used to infer their gas mass, and the temperature required to have a cloud in hydrostatic equilibrium (see \\citealt{Corbelli01}). Using a typical column density $N_H\\sim 3\\times 10^{20}\\cm^{-2}$ for a clump, and a radius $r_{ss} \\simeq 3.6 \\kpc$, a gas mass of $ M_g \\simeq 1.6 \\times 10^8 \\msun$ is inferred for the self-shielded region. In the absence of dark matter, the temperature in hydrostatic equilibrium is $T \\simeq G M_g m_H / k r_{ss} \\simeq 2.3 \\times 10^4$K. Since this is near the expected temperature in photoionization equilibrium, this would imply that the gas clumps in damped systems do not contain a lot of dark matter. However, this conclusion is altered if the gas temperature is higher due to shock heating or turbulent motions are important."
    },
    "0201/astro-ph0201043_arXiv.txt": {
        "abstract": " ",
        "introduction": "In the standard Big Bang model (SBB) the temperature of the relic radiation  from the hot phase of the Universe is predicted to increase linearly with redshift $z$: $T_{\\rm CMBR}(z) = T_{\\rm CMBR}(0)\\,(1+z)$. At the present epoch direct measurements show that $T_{\\rm CMBR}(0) = 2.725\\pm0.001$~K (1 $\\sigma$ c.l.), and that the relic radiation follows a Planck spectrum with very high precision \\cite{Mather}. However, at earlier cosmological epochs $T_{\\rm CMBR}$ cannot be measured directly. It was suggested that this scaling in proportion to $(1+z)$ can be tested by observing the population of excited fine-structure lines in the QSO absorption spectra \\cite{BW}.  The relative population of the fine-structure levels may not, however, be caused by photo-absorption of the CMBR only. Non-cosmological sources such as particle collisions, pumping by UV radiation or by IR dust emission may compete with the CMBR to populate the excited fine-structure levels. Only independent knowledge of the ambient radiation field, particle densities and the kinetic temperature of the gas allow to disentangle the  contribution of the background radiation from that of other mechanisms. For these reasons  previous studies set only upper limits to $T_{\\rm CMBR}$ \\cite{Meyer}, \\cite{Songaila}, \\cite{Lu}, \\cite{Ge}, \\cite{Roth}, \\cite{Ge2}.  The physical parameters in question can be accurately estimated if the absorber is a diffuse molecular cloud. Molecules allow to measure the volumetric gas density $n_{\\rm H}$, the kinetic temperature $T_{\\rm kin}$ and the intensity of the UV radiation field through the analysis of their distribution on the low rotational levels. In particular, intervening molecular clouds showing H$_2$ and the fine-structure lines of C\\,{\\sc i} and C\\,{\\sc ii} provides a unique opportunity to measure the cosmic microwave background radiation temperature in early cosmological epochs and to test the SBB predictions.  Another observational test of the SBB model is the measurement of the hydrogen isotopic ratio at high $z$. The standard (homogeneous) Big Bang nucleosynthesis (SBBN) predicts the same D/H abundance ratio for any direction in the early Universe since ``no realistic astrophysical process other than the Big Bang could produce significant D'' \\cite{Schramm}. Deuterium is created exclusively in BBN and therefore we can expect that the D/H ratio decreases with cosmic time due to conversion of D into $^3$He and heavier elements in stars. It is clear that the precise measurements of the D/H values at high redshift are extremely important to probe whether BBN was homogeneous. The choice of the appropriate BBN model may in turn place constrains on different models of structure formation.  In this contribution we discuss the role of the H$_2$-bearing cloud with respect to the measurement of D/H and $T_{\\rm CMBR}$ at $z_{\\rm abs} = 3.025$ toward the quasar 0347--3819. ",
        "conclusions": ""
    },
    "0201/astro-ph0201333_arXiv.txt": {
        "abstract": "We review the evolution and nucleosynthesis of AGB stars. We then discuss some of the contentious issues and quantitative uncertainties in current models. ",
        "introduction": "\\label{intro} The Asymptotic Giant Branch (AGB) phase is very short but its importance is seen in its nucleosynthesis. A revolution in stellar modelling has taken place in the last 20 years, inspired jointly by this rich nucleosynthesis and partly by new data. For example, the isotopic data coming from pre-solar grains (see this volume) forces theorists to include species that were previously ignored, species which are energetically of no importance (i.e. they play no role in determining the stellar structure) but which can be used to constrain the models. Nucleosynthesis is now important as a tracer of temperature and mixing, and not simply a by-product of energy generation. But along with these advances come more quantitative demands. It is now increasingly important to know what is known well and what is less sure. This is the goal of this paper. ",
        "conclusions": "An AGB star is a complicated thing, but that is what makes it interesting. The interplay of different kinds of physics is fascinating. Reliable quantitative estimates of their behaviour require addressing the areas of uncertainty we mention above. Before placing too much faith in these estimates, look carefully to see which assumptions went into the details-there lies the devil, as always."
    },
    "0201/astro-ph0201105_arXiv.txt": {
        "abstract": "{ We present X-ray spectral fits to a recently obtained {\\em Chandra} grating spectrum of $\\eta$~Carinae, one of the most massive and powerful stars in the Galaxy and which is strongly suspected to be a colliding wind binary system. Hydrodynamic models of colliding winds are used to generate synthetic X-ray spectra for a range of mass-loss rates and wind velocities. They are then fitted against newly acquired {\\em Chandra} grating data. We find that due to the low velocity of the primary wind ($\\approx 500 \\kmps$), most of the observed X-ray emission appears to arise from the shocked wind of the companion star. We use the duration of the lightcurve minimum to fix the wind momentum ratio at $\\eta = 0.2$. We are then able to obtain a good fit to the data by varying the mass-loss rate of the companion and the terminal velocity of its wind. We find that $\\Mdot_{2} \\approx 10^{-5} \\;\\Msolpyr$ and $v_{\\infty_{2}} \\approx 3000 \\; \\kmps$. With observationally determined values of $\\approx 500-700 \\; \\kmps$ for the velocity of the primary wind, our fit implies a primary mass-loss rate of $\\Mdot_{1} \\approx 2.5 \\times 10^{-4} \\;\\Msolpyr$. This value is smaller than commonly inferred, although we note that a lower mass-loss rate can reduce some of the problems noted by Hillier \\etal (\\cite{HDIG2001}) when a value as high as $10^{-3} \\;\\Msolpyr$ is used. The wind parameters of the companion are indicative of a massive star which may or may not be evolved. The line strengths appear to show slightly sub-solar abundances, although this needs further confirmation. Based on the over-estimation of the X-ray line strengths in our model, and re-interpretation of the HST/FOS results, it appears that the homunculus nebula was produced by the primary star. ",
        "introduction": "\\label{sec:intro} The superluminous star $\\eta$~Carinae (HD~93308, HR~4210) continues to be extensively studied over a host of different wavelengths, yet remains intriguingly enigmatic. It is amongst the most unstable stars known. In the 1840s, and again in the 1890s, it underwent a series of giant outbursts (\\eg Viotti \\cite{V1995}) which ejected large masses of material into the surrounding medium. {\\em HST} images of the resulting bipolar nebula (\\eg Morse \\etal \\cite{MDB1998}), known as the Homunculus, show it to be amongst the most spectacular in our Galaxy. The central object is now largely obscured by dust, and the cause of the outbursts and the nature of the underlying star (at outburst and today) remain speculative. The source continues to show brightness fluctuations and emission-line variations. Further details of $\\eta$~Car can be found in the review by Davidson \\& Humphreys (\\cite{DH1997}).  In recent years evidence for binarity in this system has been accumulating. Damineli (\\cite{D1996}) first noted a 5.5~yr period in the variability of the He I 10830~${\\rm \\AA}$ line. Further photometric and radial velocity studies (Damineli \\etal \\cite{DCL1997}, \\cite{D2000}), X-ray observations (Tsuboi \\etal \\cite{TKSP1997}; Corcoran \\etal \\cite{CFPSD2000} and references therein), and radio {\\mbox data} (Duncan \\etal \\cite{D1995}, \\cite{DWRL1999}) have supported the 5.5~yr period and the binary hypothesis. However, the ground based radial velocity curve was not confirmed by higher resolution spectra with STIS, indicating that at least the time of periastron passage is not well defined by the UV and optical spectra (Corcoran \\etal \\cite{CISP2001}). A comparison of the abundances from the central object(s) and the composition of the Homunculus nebula has also determined that there are at least two stars in this system (Lamers \\etal \\cite{LLPW1998}).  $\\eta$~Car is often classified as a luminous blue variable (LBV). These are massive stars believed to be in a rapid and unstable evolutionary phase in which many solar masses of material are ejected into the interstellar medium over a relatively short period of time ($\\approx 10^{4}$~yr). LBVs are regarded as a key phase in the evolution of massive stars, during which a transition into a Wolf-Rayet star occurs (\\eg Langer \\etal \\cite{L1994}; Maeder \\& Meynet \\cite{MM2001}). Due to their rarity and complex nature however, we unfortunately still have no definitive theory for mass-loss during the LBV stage. The majority of proposed mechanisms to drive LBV instabilities, the onset of higher mass-loss rates and underlying eruptions, are concerned with the importance of radiation pressure within the outer envelope of the LBV, and for example utilize pulsational instabilities (\\eg Guzik \\etal \\cite{GCD1999}), dynamical instabilities (\\eg Stothers \\& Chin \\cite{SC1993}), or presuppose Eddington-like instabilities. The latter could arise from an enhancement in opacity as the star moves to lower temperatures (\\eg Lamers \\cite{Lam1997}), or from the influence of rotation (\\eg Langer \\cite{Lan1997}; Zethson \\etal \\cite{Z1999}). Alternatively, the possibility that binarity plays a fundamental role in explaining observed LBV outburst properties has also been considered (Gallagher \\cite{G1989}), though most LBVs are not known binaries. Clearly, determining the wind and stellar properties of LBV stars is paramount (see, for example, the discussions in Leitherer \\etal \\cite{L1994} and Nota \\etal \\cite{N1996}).  An important question is the degree to which binarity influences the properties of LBVs (\\ie do LBVs in binaries evolve differently than single LBVs?). So while the presence of a companion can be exploited to help measure the mass of such stars, we must bear in mind that {\\mbox binary} LBVs and single LBVs may be quite distinct objects. Therefore, in order to use $\\eta$~Car to understand some of the defining LBV characteristics such as their extremely high mass-loss rates, we first need to determine beyond all doubt that $\\eta$~Car is in fact a binary, and then to determine the influence of the companion on the system.  Investigations over the last few years have already helped to form a basic picture of $\\eta$~Car. The orbital {\\mbox parameters}, although uncertain, indicate the presence of an early-type companion star, which will also have a powerful stellar wind. In such binaries, a region of hot shocked gas with temperatures in excess of 10~million~K is created where the stellar winds collide (Prilutskii \\& Usov \\cite{PU1976}; Cherepashchuk \\cite{C1976}). The wind-wind collision (WWC) region is expected to contribute to the observed emission from this system, particularly at X-ray and radio wavelengths. Previous X-ray observations revealed extended soft emission from the nebula and strong, hard, highly absorbed, and variable emission closest to the star (Corcoran \\etal \\cite{C1995}; Weis \\etal \\cite{WDB2001}), in contrast to the emission characteristics from single stars, which are typically softer, much less absorbed, substantially weaker and relatively constant in intensity. Since 1996 Feb $\\eta$~Car has been continuously monitored by {\\it RXTE} in the 2--10~keV band (\\eg Corcoran \\etal \\cite{CISP2001}). The lightcurve (Fig.~\\ref{fig:eta_lc}) {\\mbox contains} remarkable detail showing a slow, almost linear, rise to maximum over a period of $\\approx 1$~yr, followed by a rapid drop to approximately 1/6 of the peak intensity for $\\approx 3$~{\\mbox months}, an almost as sharp rise to approximately 1/2 of the peak intensity level, and then almost constant intensity for $\\approx 3/4$ of the proposed 5.5~yr orbital period. The drop to minimum was successfully predicted from {\\mbox numerical} models of the WWC (Pittard \\etal \\cite{PSCI1998}) before being actually observed.  \\begin{figure*} \\begin{center} \\psfig{figure=MS1953f1.eps,width=12.0cm} \\end{center} \\caption[]{Lightcurve of $\\eta$~Car observed with the {\\it RXTE} satellite and phased to the 5.5~yr orbital period. Plotted are counts detected in layer 1 of the second proportional counter unit (PCU2) and a predicted lightcurve (Pittard \\etal \\cite{PSCI1998}) from a numerical model of the wind-wind collision. The two agree well, particularly the duration of the minimum. The rise from minimum is not in good agreement, but this is thought to be due to the limitations of modelling the wind collision in 2D. The rapid change in position angle of the stars through periastron passage skews the shock cone which causes the line of sight in fully 3D models to remain in the denser wind of the primary until later phases, increasing the absorption at these times (Pittard \\cite{P2000})}. \\label{fig:eta_lc} \\end{figure*}  Small scale quasi-periodic outbursts in the X-ray lightcurve have also been detected (Corcoran \\etal \\cite{CISDPS1997}). Estimates of changes in the timescale between successive flares as a function of phase were made by Davidson \\etal (\\cite{DIC1998}) for a variety of assumed orbital elements. {\\em RXTE} X-ray observations obtained after the X-ray minimum seem to show a lengthening of the flare timescale (Ishibashi \\etal \\cite{I1999}), which indirectly support the {\\mbox binary} model. The latest published X-ray observation of $\\eta$~Car is of a high resolution grating spectrum taken with the {\\it Chandra} X-ray observatory (Corcoran \\etal \\cite{C2001}). Preliminary analysis has revealed the presence of strong forbidden line emission, which suggests that the density of the hot gas, $n_{\\rm e} < 10^{14} \\cm^{-3}$. This can be contrasted with the newly published X-ray grating spectra of the single stars $\\theta^{1}$~Ori~C (Schulz \\etal \\cite{SCHL2000}), $\\zeta$~Ori (Waldron \\& Cassinelli \\cite{WC2001}), and $\\zeta$~Pup (Kahn \\etal \\cite{K2001}), all of which have weak forbidden lines (indicative of either high densities or high UV flux near the line forming region). This lends further support to a wind-wind collision model, although it is possible that some of the forbidden emission may be related to the surrounding nebula.  If $\\eta$~Car is in fact a binary system, the orbital {\\mbox elements} and stellar parameters are not yet tightly constrained by current observations. Ground-based observations (for which good phase coverage exist) are hampered by poor spatial resolution and thus suffer contamination from strong nebular emission.  High-spatial resolution spectra have been obtained by STIS but phase coverage is currently very limited. {\\em HST} STIS observations at two different phases of the 5.52 year cycle (Davidson \\etal \\cite{DI2000}) did {\\em not} confirm the predicted variations in the radial {\\mbox velocity} of the emission lines based on the ground-based radial {\\mbox velocity} curve (Damineli \\etal \\cite{D2000}). If $\\eta$~Car is a binary, it is vitally important to determine the stellar {\\mbox parameters} of the companion so that the effect of the companion on observations can be understood and the correct stellar parameters of the primary can be derived.  \\subsection{X-ray Emission from Colliding Winds}  The wealth of information contained in X-ray spectra of colliding wind binaries (\\eg the density, temperature, {\\mbox velocity}, abundance, and distribution of the shocked gas in the wind collision region) has been a strong motivating force for observers and theorists alike in this field. Since the hot plasma in most colliding wind shocks is optically thin and collisionally ionized, and is generally assumed to be in collisional equilibrium, Raymond-Smith (Raymond \\& Smith \\cite{RS1977}) or MEKAL (Mewe \\etal \\cite{MKL1995}) spectral models are normally fitted to such data (\\eg Zhekov \\& Skinner \\cite{ZS2000}; Rauw \\etal \\cite{RSPC2000}; Corcoran \\etal \\cite{C2001}). However, the multi-temperature, multi-density nature of the WWC region means that at best simple fits with one- or two-temperature Raymond-Smith models can {\\mbox only} {\\em characterize} the broad properties of the emission. In this way one can estimate an `average' temperature of the shocked gas, and an `average' absorbing column, but little is learned of the underlying stellar wind parameters. At worst the application of one- or two-temperature models to what is inherently multi-temperature emission can lead to spurious values of some of the fit parameters, \\eg abundances (\\cf Strickland \\& Stevens \\cite{SS1998}).  Complex numerical hydrodynamical models have often been applied to gain insight into colliding wind {\\mbox systems} (\\eg Stevens \\etal \\cite{SBP1992}, Pittard \\etal \\cite{PSCI1998}). However, while undoubtedly useful, their interpretation can be {\\mbox difficult}, and to date there have been only two published papers where observed spectra are {\\em directly fitted} with synthetic spectra from such models. In the pioneering work of Stevens \\etal (\\cite{S1996}), medium resolution {\\it ASCA} spectra of the Wolf-Rayet binary $\\gamma^{2}$ Velorum were fitted against a {\\em grid} of synthetic spectra. In this fashion they were able to obtain direct estimates of the mass-loss rates and terminal velocities of the individual stellar winds. As mass-loss rates obtained from measures of radio flux or spectral line fits depend on a variety of untested assumptions, the importance of a new independent method to complement estimates from free-free radio or sub-mm observations, or from H$\\alpha$ or UV spectral line fitting, cannot be stressed enough. Rates derived by Stevens \\etal (\\cite{S1996}) with this new method were significantly lower than the commonly accepted estimates for $\\gamma^{2}$ Velorum based on radio observations, but an indication of the future benefits of this method was realized when both sets of estimates were {\\mbox later} brought into agreement following a surprisingly large reduction in the distance to this star from {\\it Hipparcos} {\\mbox data}\\footnote{The thermal radio flux, $S_{\\nu} \\propto \\Mdot^{4/3} D^{-2}$ (where $D$ is the distance to the source) whereas the X-ray flux from an adiabatic wind collision is $F_{\\nu} \\propto \\Mdot^{2} D^{-2}$. Therefore $\\Mdot_{\\rm radio} \\propto D^{3/2}$ whereas $\\Mdot_{\\rm xray} \\propto D$. If $D$ is revised downwards, $\\Mdot_{\\rm radio}$ decreases faster than $\\Mdot_{\\rm xray}$.}. We note that this method can also provide insights into the values of parameters which are otherwise virtually impossible to estimate, such as the mass-loss rate of the companion, $\\Mdot_{2}$, or the characteristic ratio of the pre-shock electron and ion temperature (Zhekov \\& Skinner \\cite{ZS2000}). The quality of recently available X-ray grating spectra now gives us access to important X-ray emission line diagnostics which should severely constrain models of the X-ray emission distribution. This means that stellar wind parameters can in principle be reliably estimated from analysis of X-ray grating spectra of colliding wind binaries.  In addition to testing the binary hypothesis, the {\\it Chandra} grating spectrum of $\\eta$~Car (Corcoran \\etal \\cite{C2001}) provides the ideal opportunity to test the method developed by Stevens \\etal (\\cite{S1996}) against a spectrum of much higher spectral resolution, and to pin down important physical parameters of the system. In this paper we therefore fit the X-ray grating spectrum using a grid of colliding wind emission models to i) test the binary hypothesis, and ii) to attempt to obtain accurate estimates of the wind {\\mbox parameters} of each star. The fact that we are able to obtain good fits, with sensible model parameters, gives us further confidence in the binary hypothesis. We also find that unlike the UV and optical work where fits to the primary are made difficult by significant contamination from the companion star, instead the X-ray emission arises from the shocked wind of the {\\em companion} and suffers essentially zero contamination from the wind of the primary. Therefore the X-ray data uniquely samples {\\mbox parameters} of the companion, in contrast to the optical analysis which probes the nature of the primary. In this sense our analysis is entirely complementary to the complex fits to the UV and optical HST spectrum of $\\eta$~Car by Hillier \\etal (\\cite{HDIG2001}). Our analysis also provides us with a new estimate of the mass-loss rate of the primary {\\mbox star}. For details of the {\\it Chandra} observation and an initial analysis of the data the reader is referred to Corcoran \\etal (\\cite{C2001}). Here we only note that there is no significant contamination of the dispersed spectrum from any spatially resolved emission (\\ie in the Homunculus). In Section~\\ref{sec:syn_spec} we discuss the creation and variation of the model spectral grid; in Section~\\ref{sec:fits} we describe the fit results; and in Section~\\ref{sec:conclusions} we summarize and conclude. ",
        "conclusions": "\\label{sec:conclusions} In this paper our aim has been to test the binary hypothesis of $\\eta$~Car by directly fitting its X-ray spectrum using a grid of spectra calculated from hydrodynamical models of the wind-wind collision. While our analysis does not prove that it is a binary, we find that the colliding wind emission model naturally provides for the range of ionization seen in the emission line grating spectrum for reasonable values of the wind parameters. We have not shown that it is inconsistent with emission from a single star, but have noted that it is unlike any of the other single stars observed so far at high energies and dispersion.  The technique applied in this paper has only been demonstrated once before and is the first time that it has been used with a high quality grating spectrum. Due to the low velocity of the primary wind ($\\approx 500 \\kmps$), most of the observed X-ray emission arises from the shocked wind of the companion star. We find it difficult therefore to fit both $\\Mdot_{1}$ and $v_{\\infty_{1}}$ as free parameters. However, the duration of the observed X-ray minimum can be used to estimate the wind momentum ratio of the stars, $\\eta$. With $\\eta$ fixed at 0.2, and $\\Mdot_{2}$, and $v_{\\infty_{2}}$ as free parameters, we are able to obtain a good fit to the data.  We find that the mass-loss rate of the companion is $\\Mdot_{2} \\approx 10^{-5} \\;\\Msolpyr$ and the terminal velocity of {\\mbox its} wind is $v_{\\infty_{2}} \\approx 3000 \\; \\kmps$. These values suggest that the companion is probably an Of supergiant (O-stars with similar wind parameters - \\eg HD~15570 (O4If), HD~93129A (O3If), HD~93250 (O3Vf), HD~151804 (O8If), and Cyg~OB2~\\#9 (O5If) - are listed in Howarth \\& Prinja \\cite{HP1989}), or is possibly a WR star. The velocity of the primary wind has been determined to lie in the range $500-700 \\; \\kmps$ (\\eg Hillier \\etal \\cite{HDIG2001}). Hence our fit implies a primary mass-loss rate of $\\Mdot_{1} \\approx 2.5 \\times 10^{-4} \\;\\Msolpyr$. From the uncertainty in the value of $\\eta$, and the interpolation on our grid, we estimate the uncertainty on our derived value for $\\Mdot_{1}$ as approximately a factor of 2.  Our best-fit estimate of $\\Mdot_{1}$ is smaller than typically inferred (\\cf Davidson \\& Humphreys \\cite{DH1997}; Hillier \\etal \\cite{HDIG2001}) However, we note that a lower mass-loss rate for the primary star can reduce some of the problems noted by Hillier \\etal (\\cite{HDIG2001}) who fitted a value as high as $10^{-3} \\;\\Msolpyr$. In particular, the models of Hillier \\etal (\\cite{HDIG2001}) suffered from absorption components that were too strong and electron-scattering wings which were overestimated. Both indicate that their chosen mass-loss rate is too high. Paradoxically both the H$\\alpha$ and H$\\beta$ {\\mbox emission} lines were weaker than observed, indicating that their mass-loss rate is too low. However, it is well known that the wind collision zone can be a strong source of emission lines (\\eg HD~5980, Moffat \\etal \\cite{M1998}), which would resolve this problem. Their inferred minimum column density is also larger than the observed X-ray value, again implying an overestimate of $\\Mdot_{1}$. To increase the mass-loss rate of the primary towards the value estimated by Hillier \\etal (\\cite{HDIG2001}), we require either a reduction in the wind velocity to $\\approx 100 \\;\\kmps$, which in turn is in conflict with previous estimates, or a reduction in the wind momentum ratio to $\\eta \\ltsimm 0.1$, which is in conflict with the observed duration of the X-ray minimum. Finally, it is worth noting that our results indicate a value for $\\Mdot_{1}$ which is closer to the value inferred for the Pistol star ($\\Mdot \\approx 4 \\times 10^{-4} \\;\\Msolpyr$; Figer \\etal \\cite{F1998}), an extreme early-type star with similarities to $\\eta$~Car.  Since current observations and theoretical modelling of the optical spectrum are unable to determine the effective temperature and the stellar radius of the primary without first determining $\\Mdot_{1}$ (\\cf Hillier \\etal \\cite{HDIG2001}), our independent estimate may prove to be extremely useful in this regard. It will be interesting to see if the results from this paper are consistent with future X-ray observations, and whether estimates of $\\Mdot_{1}$ from observations at X-ray and other wavelengths can be reconciled. Future X-ray grating observations should also help us to fix the value of the wind momentum ratio more accurately.  As the secondary wind dominates the X-ray spectrum, and its terminal velocity appears to be high ($\\approx 3000\\;\\kmps$), we should expect to see signs of Doppler broadening and shifts in the line profiles. While there is little evidence for this in the current spectrum, other orbital phases may be more favourable in this regard. The addition of Doppler effects has already been incorporated in modelled X-ray spectra (Pittard \\etal in preparation), and should provide further information on wind velocities and the structure of the wind-wind collision region.  The over-prediction of the X-ray lines in our models perhaps indicates that the companion has sub-solar abundances, which favours an O-type over a WR classification, although we would need to perform a more detailed analysis to confirm this possibility. As the primary has slightly enhanced abundances of C and N compared to solar, this suggests that to date there has been no mass exchange between the stars. Lamers \\etal (\\cite{LLPW1998}) suggested that the star which dominates the UV GHRS spectrum is not the star which ejected the nebula since the abundances in the GHRS spectrum are not as evolved as the abundances in the nebula (which are indicative of CNO-cycle {\\mbox products}). As the UV bright source is probably the companion (Hillier \\etal \\cite{HDIG2001}) this indicates that it was the primary which ejected the nebula. There are also some caveats about the analysis in Lamers \\etal (\\cite{LLPW1998}) since strong C lines (which Lamers \\etal took to indicate normal CNO abundances in the GHRS spectrum) also appear in stars known to be deficient in C (see the discussion in Hillier \\etal \\cite{HDIG2001}). It is also interesting to note that the UV spectrum brightened in 1999.1~vs.~1998.2, which is suggestive of an eclipse of the UV source near periastron in 1998. In conclusion, the high energy photons (UV and X-ray) seem to be telling us about the companion.  We emphasize that contrary to the vast majority of colliding wind systems, our X-ray analysis of $\\eta$~Car primarily probes the conditions of the shocked wind of the companion. X-ray observations of $\\eta$~Car are therefore unique in this regard since at other wavelengths (with the possible exception of the far UV) the wind of the primary dominates the observed phenomena. While our analysis has for the first time provided a direct estimate of the wind parameters of the companion star, relating these to the stellar parameters (mass, radius, luminosity) of the companion star requires more work. It is anticipated that the continued multi-wavelength study of $\\eta$~Car through and beyond the next periastron passage will further reveal the hidden secrets of this most enigmatic system.  \\newpage"
    },
    "0201/astro-ph0201349_arXiv.txt": {
        "abstract": "\\noindent \\object{RBS797} and \\object{CL 0939+4713} are two intermediate red-shift clusters ($z=0.35-0.41$). They have very different morphologies but both show surprisingly interesting structures. \\object{RBS797} looks relaxed, with an almost circular morphology; a CHANDRA observation of this cluster has revealed two deep depressions in the X-ray emission near the core. \\object{CL 0939+4713 } has instead an irregular morphology with evident substructures which seem to be in the process of merging.\\\\ Throughout this talk, if not otherwise stated, the errors are $90\\%$ confidence level.\\\\ ",
        "introduction": "The latest results from the analysis of a CHANDRA observation of the hot ($T=7.7$~keV), distant ($z=0.35$) galaxy cluster \\object{RBS797} will first be presented (see \\S~\\ref{edefilippis-B3_sec:RBS797}). The most striking observed features are two deep depressions in the X-ray emission in the core of the cluster. This is the first time such depressions have been observed in a \"distant\" cluster. The low temperature of the high density regions surrounding the holes is a clear indication that these cannot be shock regions. It is likely that the intra-cluster gas has been subsonically pushed away from the areas of low X-ray emission to the areas of higher emission by the pressure of the relativistic particles in radio lobes.\\\\ \\noindent The first results from a recent XMM observation of the galaxy cluster (\\object{CL 0939+4713} - \\object{A851}) will then be presented (see \\S~\\ref{edefilippis-B3_sec:Cl0939}). \\object{CL 0939+4713} lies at a slightly larger distance than \\object{RBS797} ($z=0.41$) but its complex and irregular structure is completely different from the almost spherical and relaxed one of \\object{RBS797}.\\\\ Its high richness, together with the presence of numerous small, blue objects in the proximity of a background quasar ($z=2.05$), has made \\object{CL 0939+4713}, in the past decade, a fundamental target in the study of the evolution of galaxy clusters.\\\\ Previous ROSAT observations have shown the cluster complex internal structure. The ROSAT analysis left, though, high uncertainties in the determination of the physical conditions of the intra-cluster gas and in the identification of its substructures.\\\\ ",
        "conclusions": "We have observed two intermediate red-shift clusters with the two new X-ray satellites: CHANDRA and XMM.\\\\ \\object{RBS797} looks almost relaxed,  \\object{CL 0939+4713} has two substructures which seem to be just in process of merging.\\\\  \\noindent With CHANDRA, thanks to its amazing angular resolution, in a recently discovered cluster such as \\object{RBS797}, we have been able to detect X-ray depressions, even in such a distant cluster; this has allowed us a straightforward prediction of the presence, in the centre of the cluster, of radio lobes generated from the central active galaxy. This prediction has been partially confirmed by low resolution radio data. New radio data, which we have now applied for, will then help us in a better understanding of the interaction between the relativistic particles and the particles in the intra-cluster gas.\\\\  \\noindent With XMM we have instead analysed a cluster which had been previously extensively studied and observed. We have resolved two main substructures forming the cluster centre and many point sources within the cluster. This is giving us the chance to perform a very accurate spectral analysis for this cluster that will allow us a study of the physical conditions of the intra-cluster gas.\\\\"
    },
    "0201/astro-ph0201455_arXiv.txt": {
        "abstract": "We analyze rotation velocities and chromospheric (\\hal) activity, derived from multi-year, high-resolution spectra, in 56 mid-M to L dwarfs.  Rotation velocities are found to increase from mid-M to L.  This is consistent with a lengthening of spin-down timescale with later type, though in the L types the trend may also be a function of stellar age.  From M5 to M8.5, a saturation-type rotation-activity relation is seen, similar to that in earlier types.  However, the saturation velocity in our case is much higher, at $\\sim$ 12 \\kms.  A sharp drop in activity is observed at $\\sim$ M9, with later types showing little or no \\hal emission, in spite of rapid rotation.  This may be due to the very high resistivities in the predominantly neutral atmospheres of these cool objects. ",
        "introduction": "In early M dwarfs (M1 - M5), the trend of lengthening spindown timescale with later type, observed in G and K dwarfs, continues: early M's are faster rotators than coeval G and K stars, and there is some evidence for increasing \\vsini from M1 to M5 (\\cite{delfosse98}).  Two mechanisms have been proposed to explain this trend:  the deepening of the convective envelope with later type (for stars with M$>$0.3\\msun, i.e., earlier than $\\sim$ M2.5 on the ZAMS; see \\cite{charb97}, and the small-scale nature of the turbulent magnetic field (for less massive, fully convective stars; see \\cite{durney93}).  The rotation-activity relation in the M1-M5 dwarfs is also similar to that in G and K stars.  A saturation-type relation is seen in coronal and chromospheric activity indicators (\\cite{delfosse98}).  Above a critical rotation velocity ($\\sim$ 5\\kms for coronal X-ray and $\\sim$ 2\\kms for chromospheric \\hal emission), activity in the early M's is saturated (at $\\sim$ 10$^{-3}$ for  $\\Lx/\\Lbol$ and $\\sim$ 10$^{-3.5}$-10$^{-4}$ for $\\Lhal/Lbol$).  Below this velocity, a range of activity levels, less than the saturation limit, are seen (though the observational sensitivity in \\vsini is too low to probe the exact trend of activity with rotation in this regime) (Fig. 1).  In stars earlier than $\\sim$ M2.5, these data are consistent with the standard rotation-activity paradigm.  In this picture, the magnetic fields that drive activity are generated by an $\\alpha\\Omega$ dynamo, whose efficiency strongly increases with faster rotation (\\cite{charb97}).  Faster rotation thus leads to more activity; saturation sets in when the fields are strong enough to cover the entire stellar surface.  One might expect the rotation-activity relation to weaken at about M2.5, since  later spectral types are fully convective.  As such, they may have only a turbulent dynamo, which is only marginally rotation-dependent (\\cite{durney93}).  However, no break in the rotation-activity relation is observed at $\\sim$ M2.5; indeed, it holds all the way down to $\\sim$ M5 (\\cite{delfosse98}).  Perhaps the turbulent dynamo becomes dominant even {\\it before} full convection sets in, and is efficient enough, even though only weakly rotation-dependent, to induce saturation at low \\vsini in these stars.  This would induce a smooth transition in the rotation-activity relation across the full-convection boundary.  In this paper, we study the \\vsini and chromospheric \\hal emission in 56 mid-M to L dwarfs (M5 - L6), all of which are field objects.  Our aim is to examine whether the trends in the rotation and rotation-activity connection noted above continue as we move beyond early M, to still later spectral types.  A preliminary version of this work was presented at Cool Stars XI by Basri (2000).  In the present work, the previous sample has been extended to include more objects, and \\vsini, \\hal flux, spectral type and \\teff are determined more accurately. ",
        "conclusions": "\\begin{itemize} \\item Rotation velocities continue to increase from mid-M to L.  In the M dwarfs, this continues the trend of lengthening spindown timescale with later type, that is seen in earlier types.  In the L dwarfs, this trend may be due to substellarity and thus comparative youth, or a true increase in spindown timescale, perhaps due to the high atmospheric resistivities. \\item A saturation-type rotation-activity connection is observed from M5-M8.5, similar to that in the early M's.  However, the critical velocity for saturation ($\\sim$ 12 \\kms) is much higher than in the early M's ($\\sim$ 2 \\kms).  No obvious correlation is seen between rotation and activity below 12 \\kms.  This, together with the high critical velocity, may be evidence for a turbulent dynamo that is only weakly dependent on rotation. \\item A sharp drop in activity is seen at $\\sim$ M9, with later types showing little or no evidence of \\hal emission, in spite of being very rapid rotators.  This may be due to the very high resistivities in the cool, predominantly neutral atmospheres of late M and L dwarfs:  the resultant rapid decay of currents damps the build-up of large magnetic stresses, which could otherwise energetically support a chromosphere and corona and drive activity. \\end{itemize}  \\plotfiddle{mohantys1.eps}{8cm}{0}{45.}{45.}{-150cm}{0cm} \\figcaption{\\label{fig1} $\\Lhal/\\Lbol$ versus \\vsini for M4 and M5 dwarfs (from \\cite{delfosse98}).  Upper limits in \\hal emission and \\vsini marked by arrows (though actual equatorial velocity may be higher) }  \\plotfiddle{mohantys2.eps}{7cm}{0}{45.}{45.}{-150cm}{0cm} \\figcaption{\\label{fig2} \\vsini versus Spectral Type, for M5 to L6 dwarfs.  Overlapping objects marked with spines.  Horizontal line is at 15 \\kms, below which we define objects as slow rotators.  Vertical line marks spectral type L1.5, above which all objects rotate rapidly.  `Li' marks objects with Lithium; these are confirmed brown dwarfs. }  \\plottwo{mohantys3a.eps}{mohantys3b.eps} \\figcaption{\\label{fig3}LEFT: $\\Lhal/\\Lbol$ versus \\vsini for M5 to L6 dwarfs.  {\\it Top Panel}:  M5 - M8.5.  M5-6.5 in red; M7-8.5 in brown.  {\\it Bottom panel}:  Our entire sample down to L6.  M9-9.5 in black; L0 and later, with \\hal emission, in light green; L0 and later, with no detected \\hal emission, in dark green (these objects are also marked with downward arrows).  Diamond marks LP 944-20, a confirmed brown dwarf that has been observed to flare in X-rays and radio.  RIGHT:  $\\Fhal$ (\\hal surface flux) versus \\vsini.  Otherwise same as left panels.  In the M5 - M8.5 dwarfs, using surface flux instead of luminosity ratio as activity indicator seems to give a tighter rotation-activity relation (the two outliers with the highest $\\Fhal$ are flare stars), but perhaps not enough to choose one over the other as a more fundamental measure of chromospheric activity.  }  \\plotfiddle{mohantys4.eps}{7cm}{0}{45.}{45.}{-150cm}{0cm} \\figcaption{\\label{fig4} $\\Fhal$ versus \\teff for M5 - L6 dwarfs.  Note rapid falloff in activity at about 2400 K ($\\sim$ M9).  Plot reveals that the decline in activity with later type may be a direct result of decreasing \\teff.  Note that the slope of the envelope of undetected $\\Fhal$, for \\teff $\\leq$ 2300 K, reflects the improvement in \\hal detection limits with decreasing \\teff (since the background photosphere becomes fainter), and not a real trend in $\\Fhal$; the actual \\hal surface flux here could be much lower. }"
    },
    "0201/astro-ph0201513_arXiv.txt": {
        "abstract": "New theoretical results from large-scale relativistic close coupling calculations reveal the precise effect of resonances in collisional excitation of {\\sc x}-ray lines of Ne-like \\fe17. Employing the Breit-Pauli R-matrix method and a 89-level eigenfunction expansion including up to $n = 4$ levels shows significant resonance enhancement of the collision strengths of forbidden and intercombination transitions. The present results differ from all previous calculations, heretofore without detailed resonance structures, and should help resolve longstanding discrepancies. In particular, the present line ratios of three benchmark diagnostic lines 3C, 3D, and 3E at 15.014, 15.265, and 15.456 $\\AA$ respectively, are in excellent agreement with two independent measurements on Electron-Beam-Ion-Traps [Laming \\etal, Astrophys.~J {\\bf 545}, L161~(2000) and Brown \\etal, Astrophys.~J {\\bf 502}, 1015~(1998)]. The strong energy dependence due to resonances in these and other cross sections is demonstrated for the first time. It is of general importance and strongly manifests itself in {\\sc x}-ray plasma diagnostics. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0201/astro-ph0201039_arXiv.txt": {
        "abstract": "Very sensitive \\hi\\ 21cm observations have been made in 860 directions at  $\\delta \\geq -43\\arcdeg$ in search of weak, Galactic,  high-velocity \\hi\\ emission lines at moderate and high Galactic latitudes. One-third of the observations were made toward extragalactic objects that are visible at optical and UV wavelengths. The median rms noise in the survey spectra is 3.4 mK, resulting in a median $4\\sigma$ detection level of $N_{HI} = 8 \\times 10^{17}$ cm$^{-2}$ averaged over the $21 \\arcmin$ beam of the telescope. High-velocity \\hi\\ emission is detected in 37\\% of the directions; about half of the lines could not have been detected in previous surveys.  The median FWHM of detected lines is 30.3 km~s$^{-1}$. High-velocity \\hi\\ lines are seen down to the sensitivity limit of the survey implying that there are likely lines at still lower values of $N_{HI}$.  The weakest lines have a kinematics and distribution on the sky similar to that of the strong lines, and thus do not appear to be a new population. Most of the emission originates from objects which are extended over several degrees; only a few  appear to be compact sources.  At least 75\\%, and possibly as many as 90\\%, of the lines are associated with one of the major high-velocity complexes.  With the increased sensitivity of this survey, the Magellanic Stream is seen to extend at least  $10\\arcdeg$ to  higher Galactic latitude than previously thought and to be more extended in longitude as well. Wright's Cloud near M33 has an extended low-$N_{HI}$ component in the direction of the Magellanic Stream. The bright \\hi\\ features which have dominated most surveys may be mere clumps within larger structures, and not independent objects. Although there are many lines with low column density, their numbers do not increase as rapidly as $N_{HI}^{-1}$, so most of the \\hi\\ mass in the high-velocity cloud phenomenon likely resides in the more prominent clouds. ",
        "introduction": "Emission in the 21 cm \\hi\\ line is seen over a large fraction of the sky at velocities $-500 \\lesssim V_{LSR} \\leq -100 $ km~s$^{-1}$ and $+100 \\leq V_{LSR} \\lesssim +400 $ km~s$^{-1}$: too large to arise solely from Galactic rotation, yet seemingly not part of the Hubble flow  (see \\citet{wvw97} for a recent review; throughout this work we use the definition that high-velocity \\hi\\ has $|V_{LSR}|\\geq100$ km~s$^{-1}$ in directions where such velocities are not expected from Galactic rotation). These  high-velocity clouds were discovered in the 21cm line of \\hi\\  \\citep{muller}, and for many decades the 21cm line was their primary spectral signature. Surveys of the sky in search of high-velocity \\hi\\ are thought to be complete at the level of  $N_{HI} \\geq 2 \\times 10^{18}$ cm$^{-2}$  for objects with an angular size of a degree or greater \\citep{wakk91}.  The most prominent southern high-velocity clouds can be associated with the Magellanic Clouds \\citep{mathewson}, and some high-velocity clouds lie in the Galactic halo \\citep{vanwoerden99}, but the origin of most of the  emission is unknown. It has been suggested that  high-velocity clouds result from phenomena as diverse as  extensions of the Galactic disk, material accelerated by supernovae or other energetic events, the remnants of a  Galactic fountain,  condensations in a Galactic halo, infall of debris from satellite galaxies, and objects in the Local Group of galaxies \\citep{oort66, habing, davies, shapiro, bregman80, norikeu, blitz, kalkerp99, mallouris}. It is probable that high-velocity clouds are a heterogeneous population and  all of these possibilities are correct for different subsets of the data.  A number of high-velocity absorption lines have been detected in the ultraviolet  toward extragalactic sources (e.g., \\citet{savage93, bowen93, bowen95, Lu94, murphy00, gibson2000}). The UV transitions  are much more sensitive to small amounts of neutral gas than the 21cm line: a cloud with $N_{HI} = 10^{18}$ cm$^{-2}$ will produce  strong Mg II absorption  even if the metallicity is only 0.1 solar \\citep{ls95,charleton00}. The statistics of the UV absorption lines suggest that high-velocity clouds are more abundant than would be expected from  21cm \\hi\\ surveys \\citep{savsem96}. The extrapolation from the 21cm measurements to lower $N_{HI}$ levels is always uncertain  \\citep{wakk91}, and the UV results are based on relatively few sight lines, but these observations raise the interesting possibility that the Milky Way is surrounded by a mist of high-velocity gas with $N_{HI} \\lesssim 10^{18}$ cm$^{-2}$ which has escaped detection in previous surveys.  This is doubly interesting because the  HI disks of other spiral galaxies are truncated at  the level of a few $10^{19}$ cm$^{-2}$, presumably because of ionization  by the extragalactic radiation field \\citep{vangorkom91, vangorkom93, bochk, maloney, corbelli}.  This possible population of clouds with low $N_{HI}$  might have a different distribution and different properties than the denser high-velocity clouds.  It might also contain significant mass.   To try to bridge the gap between the sensitivities of the 21cm and the UV measurements, we began in 1992 a very sensitive 21cm \\hi\\  search for high-velocity lines too faint to have been detected in previous 21cm surveys.  Some results from the first third of the data have already been published \\citep{mls95}. Since  the beginning of this project there has been renewed interest in the high-velocity cloud phenomenon. There has been an extensive 21cm survey of the southern high-velocity cloud population \\citep{putmanetal2002}. There has been a proposal that most of the clouds are members of the Local Group at a distance $\\sim 1$ Mpc \\citep{blitz}, and another that  some fraction of them are dwarf irregular galaxies which lack bright stars \\citep{bb99,bb00}. There have been investigations into the connection between high-velocity clouds and the neutral gas in the local universe \\citep{charleton00, zwaan00, zwann01}. While the data presented here bear somewhat on these issues, the main goal of this work has always been to investigate the high-velocity cloud phenomenon  at low values of $N_{HI}$, and that aspect of the data is the focus of this paper. ",
        "conclusions": "The high-velocity \\hi\\ emission detected in this survey covers 37\\% of the sky, and lines are found down to  the median completeness limit of $8 \\times 10^{17}$ cm$^{-2}$ and below. There is a nearly constant covering fraction per unit log($N_{HI}$) from $N_{HI} = 10^{19}$ cm$^{-2}$ down to the limit of the survey, and there is no indication  of a cutoff in the high-velocity \\hi\\ population at low $N_{HI}$. Extrapolating this trend  to still lower $N_{HI}$ implies  a covering fraction near $60\\%$ for $N_{HI} \\geq 10^{17}$ cm$^{-2}$.  It is clear that the limiting $N_{HI}$ of the high-velocity cloud phenomenon has yet to be established.  High-velocity \\hi\\ emission lines can be characterized as either ``wings'' or ``Gaussian components'', the difference being whether they are blended with lower-velocity emission and thus have uncertain properties (the wings), or can be fit by a  Gaussian function whose parameters are likely to be meaningful (the components).  There are 309 detected Gaussian components, found in 28\\% of the survey directions. Much of our analysis concentrates on the components (Table 4) because they have accurate velocities and their values of  $N_{HI}$ do not depend critically on the precise velocity definition of a ``high-velocity'' line.  However, the Gaussian components are only part of the high-velocity \\hi\\ phenomenon.  The vast majority of the weak lines detected in the survey seem related both spatially and kinematically to the lines known from previous work. Weak lines are not found outside the $V_{LSR}$ range of the strong lines, and there is no evidence that the weak lines merge into the Hubble flow. The integrated emission of the high-velocity components is skewed to negative LSR velocities but has a mean near zero when motion of the LSR about the Galactic Center is removed.  Most of the lines detected in this survey, at all values of  $N_{HI}$, are associated with complexes like  the Magellanic Stream and Complex C.  The fraction of association may be as high as $90\\%$. At the sensitivity of the survey the Magellanic Stream appears to extend at least $10\\arcdeg$ to higher Galactic latitude than previously thought, and, near its tip, broadens in longitude as well.  Low-$N_{HI}$ lines associated with known complexes can be found far outside their traditional boundaries, e.g., in the Magellanic Stream and Wright's cloud. The covering fraction of the sky increases as lower $N_{HI}$ lines are measured mainly because the size of individual clouds grows. Our data thus suggest that, at the sensitivity level of the survey, most high-velocity clouds do not have sharp edges in \\hi.  Compact high-velocity clouds have received recent scrutiny because they may be at a larger distance than the principle complexes \\citep{bb99, bbc01}. There are a few low-$N_{HI}$ lines which appear spatially isolated in the survey data (Table 9), but the majority of the emission appears to be extended, and in some directions  it is extended at low $N_{HI}$.  It would be quite valuable to observe  these with higher angular resolution to see if the material is clumped or remains fairly smooth.  Despite the abundance of low-$N_{HI}$ lines, their numbers are not sufficient to contribute substantially to the mass of high-velocity \\hi\\ unless, against the evidence, they are located far beyond the bright clouds. This was the conclusion also of \\citet{colgan}, and \\citet{mls95}. In a  set of compact, isolated, high-velocity clouds mapped at high angular resolution, the extended halos  contain most of the \\hi\\ flux \\citep{bbc01}.  In this respect the compact clouds may have a different structure than the more general high-velocity emission.  In contrast to spiral galaxies, whose \\hi\\ disks have sharp edges and do not increase in size as observational sensitivity is increased \\citep{vangorkom93}, the area covered by high-velocity \\hi\\ clouds is a strong function of the sensitivity of the observations (cf the current results with those of \\citet{giovanelli}, \\citet{wakk91}, and \\citet{colgan}). And yet, the high-velocity emission seems so concentrated into complexes that observation in only a few directions may easily produce misleading impressions about its global properties (e.g., \\citet{bowen95}).  For this reason we believe that many more sight-lines must be studied in the UV before it can be determined if there is a statistical excess in the occurrence of, e.g., Mg II high-velocity lines compared to 21cm high-velocity lines.  If the bright \\hi\\ lines which constitute the ``classic'' high-velocity cloud population are merely the peaks in a more extended medium, then the validity of statistical analyses based on cloud counts must be called into question.  A clumpy cloud may be counted several times if observations are not sensitive enough to reveal the gas between the clumps.  Until there is a better way of specifying the boundaries of a cloud than a noise level set by instrumentation, we believe that statistical analyses of cloud counts are in danger of giving  misleading results.  There are  high-velocity clouds which are predominantly ionized  at $\\sim 10^4$ K, there are clouds with a significant component of mass at $T \\sim 3 \\times 10^5$ K, and there may even be clouds which are sources of soft X-ray emission \\citep{sembach95, weiner96,  tufte98, kerp99, sembach99,sembach00}.  This suggests that the observed  low-$N_{HI}$ lines in some directions might arise in  a trace neutral component of  an otherwise ionized cloud. The gas toward III Zw 2, where there is high-velocity Mg II absorption ($\\S6.1$) but no detectable 21cm  \\hi\\ line,  may be in a similar state. These considerations reenforce the view  that bright \\hi\\ clumps are merely one aspect of the high-velocity cloud phenomenon."
    },
    "0201/astro-ph0201507_arXiv.txt": {
        "abstract": "The {\\em Infrared Space Observatory} observed over 900 objects with the Short Wavelength Spectrometer in full-grating-scan mode (2.4--45.2 \\mum). We have developed a comprehensive system of spectral classification using these data.  Sources are assigned to groups based on the overall shape of the spectral energy distribution (SED). The groups include  naked stars, dusty stars, warm dust shells, cool dust shells, very red sources, and sources with emission lines but no detected continuum. These groups are further divided into subgroups based on spectral features that shape the SED such as silicate or carbon-rich dust emission, silicate absorption, ice absorption, and fine-structure or recombination lines. Caveats regarding the data and data reduction, and biases intrinsic to the database, are discussed. We also examine how the subgroups relate to the evolution of sources to and from the main sequence and how this classification scheme relates to previous systems. ",
        "introduction": "Spectral classification organizes astronomical sources into groups with similar properties based on the general or detailed morphology of their spectral energy distributions (SEDs). Consequently, the classification criteria depend on the wavelength region and spectral resolution used. Both of these parameters must be uniform in order to create consistent criteria for arranging the sources in a database. The similarities and differences that result from applying a successful classification system to a sufficiently large sample of sources not only improve our knowledge about the sources but provide a basis for understanding the physical parameters of the objects.  The best example of how a classification system can lead to insight into the physical properties of the objects studied is provided by optical spectral classification \\citep[cf.][]{hs86}. From the earliest systems based on general color \\citep[e.g.][]{ruth63}, several competing systems emerged based on spectral line ratios \\citep[e.g.][]{sec66,sec68,vog74,vw99, pic90}.  Of these, the Harvard system used originally in the Draper Memorial Catalogue \\citep{pic90} grew to predominate due to the large numbers of sources classified ($>10,000$), and served as the basis for the Henry Draper Catalogue \\citep[beginning with][]{cp18}.  The MK spectral classification system evolved from the Harvard system \\citep[e.g.][]{mor38,mkk43}. This two-dimensional system provided the clues necessary to disentangle the different stages of the life cycle of a star and the relation of intrinsic parameters such as mass and metallicity to directly observable properties.  MK spectral classification remains the single most powerful diagnostic tool available to astronomers when applied to naked stars, i.e. stars not embedded in dust.   Unfortunately, the very early and very late stages of stellar evolution rarely involve naked stars. The sources are deeply embedded within interstellar dust clouds or circumstellar dust shells, either of which absorb the optical radiation and re-emit it in the infrared.  This dust can absorb so much of the optical radiation from the star that traditional classification based on the photospheric properties of the star in the optical is difficult, if not impossible.  Near-infrared observations can often penetrate the obscuring dust, permitting direct measurements of the stellar photosphere.  The spectral region between 1 and 9 \\mum\\ is rich in atomic and molecular lines which trace temperature and luminosity.  For example, CO, SiO, and water vapor are sensitive indicators in oxygen-rich stars, even with low spectral resolution; the Phillips and Ballick-Ramsey C$_2$ bands as well as CN and CO serve for carbon stars.  However, the emission from the dust distorts the photospheric continuum and fills in the absorption features, making analysis difficult. Observations in the thermal infrared trace the emission from the dust itself.  The characteristic SED of the dust is distinctive enough to serve as the basis for classification \\citep{lmp86,lm87,css89}.   The infrared spectra obtained by the Low-Resolution Spectrometer (LRS) on the {\\em Infrared Astronomical Satellite} (\\iras) are the best example of a nearly complete, self-consistent database that is ideal for spectral classification. These spectra cover wavelengths from 7.7 to 22.7~\\mum\\ at a spectral resolution of $\\lambda/\\Delta\\lambda\\sim$20--60. The original LRS atlas contained  spectra from 5,425 sources \\citep{lrs86}. \\cite{vk91} expanded the database to 6,267 and  \\cite{kvb97} extracted almost 5000 additional spectra from the raw data to create a spectral database of 11,224 sources, making the LRS observations the largest infrared spectral database to date.  This database includes most of the 12 \\mum\\ objects in the sky brighter than 10 Jy at 12 \\mum\\ (magnitude +1),  and several infrared classification systems have been developed from it.  The initial LRS classification scheme \\citep{lrs86,lrs88} sorted the original database of 5,425 sources into 10 groups, essentially  based on the dominant spectral feature in the 10 \\mum\\ region. These groups were subdivided further, usually by the strength of the dominant feature.  The AutoClass algorithm (also known as AI for artificial intelligence) used a Bayesian algorithm to sort the database into self-consistent classes with no a priori input about the nature of the spectra \\citep{css89, gsv89}. \\cite{kvb97} used one-letter codes to identify the character of each spectrum in the expanded database (11,224 sources).  These various classification systems have divided the LRS database into distinct sets of spectral classes. However, none of these  systems has been applied to a substantial number of spectra from instruments other than the LRS.  Other schemes focused on subsets of the LRS database.  For example, \\citet[][hereafter, LML]{lml88,lml90} classified evolved oxygen-rich stars based on their dust emission characteristics.  This system, as modified by \\citet[][hereafter SP]{sp95,sp98}, has also been applied to ground-based spectral measurements \\citep[e.g.][]{ce97, mgd98}.  Spectra  taken by the Short Wavelength Spectrometer (SWS) on the {\\em Infrared Space Observatory} \\citep[\\iso][]{kes96,deg96} are now publically available. In this paper, we focus on the full-range, moderate-resolution spectra obtained in the SWS01 observing mode. These observations are over a greater wavelength range than the LRS database (2.4--45.2 \\mum\\ compared to 7.7--22.7 \\mum), and at a higher spectral resolution ($>$300--400 vs. 20--60).  The SWS01 spectral resolution is sufficient for detailed examination of band structure and atomic fine-structure lines.  The extended wavelength range includes both the near-infrared spectral region, which is dominated by molecular bands from stellar photospheres, and the thermal infrared region, which is dominated by dust emission.  The LRS database is compromised by inadequate wavelength coverage on the short-wavelength side of the strong spectral features produced by silicate dust (10~\\mum) and silicon-carbide grains (11.5~\\mum), making it difficult unambiguously define the stellar continuum.  \\iso\\ obtained observatory-style pointed observations whereas \\iras\\ obtained spectra as an adjunct to the main survey with the LRS as a secondary instrument.  Consequently, the SWS database only contains full-range spectra of $\\sim910$ specifically targeted sources (1248 total spectra, including duplicates and off-positions).  To ensure that \\iso\\ obtained SWS spectra of as wide a variety of sources as possible, the observing lists of the STARTYPE proposals\\footnote{The STARTYPE proposals received \\iso\\ project names STARTYP1, STARTYP2, and ZZSTARTY.} targeted sources in categories which were under-represented in the infrared classification systems (\\S \\ref{sec.sample}).  The result is a robust database of infrared spectra which  is the basis for our infrared spectral classification system.  We describe the sample of the observed sources and the structure and calibration of the spectral data in Section 2.  Section 3 details the criteria for the classification system, which we discuss in Section 4. The actual classifications are presented in Appendix A. ",
        "conclusions": "}   \\subsection{Calibration Issues and the Classifications \\label{sec.caliss}}  As mentioned in \\S \\ref{sec.data}, the browse products used to classify most of the spectra did not fully correct for flux discontinuities between bands.  The most challenging normalization problems occur between Bands 2C and 3A, at $\\lambda\\sim12$~\\mum, and between Bands 3D, 3E, and 4, at $\\lambda\\sim$26--30~\\mum.   We discuss them briefly here to the extent that they influence the classification effort.  Spectra in Group 2 are most sensitive to discontinuities and memory effects near 12~\\mum, because the shape of the emission and absorption features in the 10--12~\\mum\\ region serve as the primary features for classification into the subgroups.  If the flux discontinuity is simply related to a gain difference between bands, normalization during reprocessing (if needed at all) would simply scale Band 3A to match 2C without changing the basic shape of any features present.  If the discontinuity results from memory effects, however, it is more problematic.  Even with the {\\it dynadark} correction and normalization some error may remain in the shape of the spectrum. Fortunately, this problem does not compromise the classification of a spectrum as oxygen or carbon-rich, although it might cause a spectrum to be (mis)classified as 2.SEa instead of 2.SEb, for instance.  Normalization of Bands 3D, 3E, and 4 is complicated by a light leak and by the unreliability of 3E. Most spectra show a smooth shape with a roughly constant slope from Band 3D through Band 3E and into Band 4, which allows a straightforward normalization of these bands to each other.  However, spectra with structure near 26~\\mum\\ present more of a problem, since the changing slope of the spectrum makes extrapolation across Band 3E difficult. This problem affects the carbon-rich sources in Groups 3 and 4 most significantly and  limits our confidence in the shape of the emission feature in the 26--30~\\mum\\ region.  In the browse product spectra produced from OLP 7.1, the normalization of the segments makes the 26--30~\\mum\\ feature appear narrow and peaked around $\\sim$25--29~\\mum.  Applying our normalization algorithm to data in OLP 10.0 broadens the feature to $\\sim$25--34~\\mum.  The literature tends to refer to this feature as the 30~\\mum\\ emission feature, possibly attributable to MgS \\citep{gm85,beg94}.  With the current uncertainties in calibration, we are unable to definitively address this issue.  To date, no model has been developed to correct the memory effects in Band 4.  The entire shape of Band 4 can be compromised, and, in terms of the spectral classification, this influences whether a spectrum is classified in Group 4 or 5.  For example, a spectrum could be misclassified, as a 4.PN instead of a 5.PN because the spectrum appears to have turned over in Band 4 when it is actually still climbing.  The memory effect in Band 4 can also influence our ability to recognize crystalline silicate features, especially at 40 and 43~\\mum.  These features could be washed out when the two scan directions are combined because of the difference in flux levels between them.  As with the other issues raised here, the Band 4 memory effect should have a limited impact on the classifications.  Despite these issues, the basic classification scheme and the grouping of the spectra should prove robust.  The movement of a few spectra from CR to CE or from Group 4 to Group 5 will not change the overall nature of the database or the existence of any of the evolutionary patterns discovered therein (\\S \\ref{sec.co}--\\ref{sec.crich}).   \\subsection{Comparison With \\iras\\ Classifications \\label{sec.irascomp}}   Although we are dealing with a non-uniform database (\\S \\ref{sec.seleff}), we can still compare our classifications with the LRS classes \\citep{lrs88} and the classes of \\citet[][hereafter KVB]{kvb97}. Only the subset of SWS sources with LRS classifications (379 sources) or KVB (567 sources) classifications can be considered, so the numbers quoted below will not be the same as those given for each subgroup in Table \\ref{tab.cross1-2b}. Also, recall that for LRS class $1n$, $n=2\\beta$ where $\\beta$ is the spectra index:  \\begin{eqnarray} F_{\\lambda} & \\propto & \\lambda^{-\\beta}. \\end{eqnarray}  \\noindent Thus, when $\\beta=4$, the spectrum behaves as a pure Rayleigh-Jeans tail and is in LRS class  18.  Sources with low-contrast dust mimic lower spectral indices and receive lower LRS characterizations.  For example, LML and SP showed that many sources in LRS classes  13--16 show low-contrast alumina-rich dust in their spectra.   \\subsubsection{Similarities \\label{sec.sim}}  Group 1, the dust-free stars, corresponds well to the LRS classes 17--19. Of the 60 objects in Group 1 with LRS classifications, 54 are in LRS classes 17--19, with 39 in class 18. Of the 2.SEa sources, with low-contrast dust, 81\\% of the 53 sources are in LRS classes 13--16, as expected. Similarly, the oxygen-rich  dust sequence, described below in \\S \\ref{sec.orich}, should begin in the  $2n$ range and progress to the $3n$ range, where $2n$ corresponds to silicate emission and $3n$ to silicate absorption. Nearly 90\\% of the 67 objects in subgroups 2.SEb, 2.SEc, and 3.SE have LRS classes $2n$. The sources in 3.SB, the self-absorbed subgroup, are split between $2n$ and $3n$, and 11 of 12 sources in 4.SA are $3n$ or $7n$ (recall that $7n$ is the red counterpart of $3n$). In the carbon-rich sources, 31 of 37 sources in subgroups 2.CE, 3.CE, and 3.CT have LRS=$4n$, the carbon-rich LRS class. Only about a third (14 of 41) of the PNe subgroups 4.PN, 4.PU, and 5.PN, have LRS classifications, but those that do tend (11 of 14) to be $9n$, that is, red objects with emission lines but no detected 11.3 \\mum\\ UIR feature. For the young, red sources in Groups 4 and 5, even fewer, $\\sim25$\\%, have LRS classifications, so small numbers make valid comparisons problematic.  Still, most of those with LRS data in our SA or SE subgroups do have silicate absorption or emission LRS classifications.   A comparison of our classifications with those of KVB shows comparable similarities. For example, 37 of the 40 sources with their class C, for carbon-rich, are in one of our carbon-rich subgroups (mostly 2.CE). More than 80\\% of their A (10 \\mum\\ absorption) sources are in our SA or SB subgroups and more than 90\\% of their E (10 \\mum\\ emission) sources are in our silicate emission subgroups. Almost 90\\% of their S (stellar) sources are in Group 1, our naked star category.  \\subsubsection{Distinctions \\label{sec.diff}}   \\begin{figure} \\plotone{f6.eps.orig} \\caption{Comparison of spectra with the same LRS class but different KSPW classes. Top: LRS=80 (UIR emission): HR 4049 2.U and HD 97048 5.U (smoothed). Middle: LRS=32 (blue SED + silicate absorption): WR 112 3.W and V645 Cyg 5.SA. Bottom: LRS=41 (C-rich): RY Dra 2.CE and IRAS 22303+5950 4.CR. } \\label{fig.lrscat} \\end{figure}   While the overall correspondence betweeen our classifications and those from LRS-based schemes is reasonable, there are a number of important differences.  For instance, misidentification of UIR features as silicate absorption occurred in the LRS classifications due to the low spectral resolution and bandwidth. This is largely avoided in the SWS database because of the higher spectral resolution and especially the expanded bandwidth. The extended wavelength coverage allows confirmation of suspected 7--11 \\mum\\ UIR features with those at 6.2 and 3.3 \\mum\\ which were outside the LRS range.   It was mentioned above that most of the 4.SA and 5.SA sources  (24 of 27) were in LRS classes corresponding to silicate absorption. However, 17 of those sources were $3n$, with ostensibly blue SEDs. Characterizations of $7n$ would have been more correct, but the short wavelength cutoff of only 7 \\mum\\ presumably prevented an accurate assessment of the overall SED. In the 5.UE group,  25 sources had red LRS characterizations, but less than half (10) were $8n$, the UIR+emission line class. Lack of sensitivity of the LRS precluded the detection of UIRs in some sources, while the limited spectral range caused the confusion of UIRs and silicate absorption in others.   Examination of the carbon-rich classes and ostensibly carbon-rich objects further illustrates  the limitations of the old LRS classifications when dealing with SWS data. One source, AFGL 2287, was classified as carbon-rich in all three LRS-based schemes but is classified by us as self-absorbed silicate emission (3.SBp). In the limited spectral range of the LRS data, self-absorbed silicates can appear similar to carbon-rich spectra \\citep{wc88}. However, with the SWS, AFGL 2287 can be seen to have none of the other features typical of carbon-rich sources such as the absorption features at 13.7 \\mum\\ or 3 \\mum.  Seven other sources were classified as carbon-rich in the LRS atlas but as oxygen-rich by AutoClass and KVB. With the high-quality data from SWS, we can confirm that they are indeed oxygen-rich\\footnote{They are ST Her (2.SEa), FI Lyr (2.SEa), Z Cas (2.SEap), $\\pi^1$ Gru (2.SEa), AD Per (2.SEa), AFGL 2199 (3.SB), and AFGL 1992 (3.SB).}.  As an example of a discrepancy in the opposite sense, only one of our 4.CR sources has a carbon-rich LRS class, while most  of the rest (8 of 11) are classed as 21--23, low-contrast silicate emission. The AI classifications also mistook most  of 4.CR (9 of 10) for oxygen-rich ($\\zeta4$) because of the limited spectral range on which the classifications were based. Of the 31 objects observed with SWS which SIMBAD lists as carbon stars, only two (FI Lyr and CIT 11) are in non-carbon KSPW subgroups; four more are in Group 7 or flagged, so 25 of 27 agree with SIMBAD. The LRS scheme, on the other hand, has only 16 of 24 sources with $4n$ or 04 designations. Again, the superior sensitivity  and spectral range of SWS enabled the proper classification of these sources as carbon-rich.   The LRS  $4n$ classes base the second digit on the strength of the 11.5 \\mum\\ silicon carbide feature. However, nothing in the LRS classification indicates the shape of the underlying continuum for carbon-rich objects. Thus, sources as dissimilar as RY Dra (2.CE) and IRAS 22303+5950 (4.CR) possess the same LRS classification 41 because both have weak 11.5 \\mum\\ features (Fig. \\ref{fig.lrscat}). LRS class 44 contains both W Ori  (2.CE) and IRC +50096 (3.CE) despite their distinctly different underlying SEDs.  There simply are {\\em no} appropriate categories in the LRS scheme in which to place {\\em any} of the carbon-rich sources in our Groups 3 and 4 without loss of significant information about the spectra.  Similarly, placing all of our Group 1 naked stars into $1n$ would also lose important information about the photospheric chemistry (1.NO vs. 1.NC) or the presence of emission lines (1.NE). Other KSPW classes with no good LRS counterparts include 2.E, 2.U, 2.C/SE, 4.SC, 4.U/SC, 4.SEC, and 6.  Other LRS classes also   mis-matched sources, as Figure \\ref{fig.lrscat} shows. These disparate sources were placed in the same LRS classes because of the limited spectral range of the LRS and small number of features considered in that classification scheme. Table \\ref{tab.irascomp} compares the KSPW classes to the LRS classes. Column 2 lists the KSPW classes that are well-matched to an LRS class as defined in the \\citet{lrs88}, for example 1.N and the $1n$ class. Column 3 lists our groups that could be placed in an LRS class, but only with information lost, such as 2.C/SE in $2n$ or 4.U/SC in $8n$. The last column lists the other KSPW groups in which the LRS classes actually appear but are not well-suited to each other. Application of the LRS scheme to classify the SWS database would have essentially  ignored the additional information gained from the larger bandwidth, higher spectral resolution, and greater sensitivity of SWS.      \\subsection{Clarifying Evolutionary Patterns}  \\subsubsection{The CO Paradigm \\label{sec.co}}  The search for patterns and relations among the infrared spectral classifications identified here must first consider the observed dichotomy between carbon-rich and oxygen-rich dust chemistry. In evolved stars, the chemistry of the dust depends on the C/O ratio of the material ejected from the envelope.  The formation of CO molecules will exhaust the less abundant of carbon or oxygen, leaving the other element available to form molecules which serve as seeds for dust formation.  This CO paradigm works admirably well.  In only a few cases does the chemistry of the stellar photosphere differ from that of the dust, and most of these cases probably arise within binary systems.  For example, the silicate carbon stars (2.C/SE) show carbon-rich photospheric features and oxygen-rich dust.  In these sources, the dust emission may arise from a disk around an unseen companion which trapped mass lost from the primary before it evolved into a carbon star \\citep[see][for a recent study of the SWS spectrum of V778 Cyg and a discussion of competing models]{yam00}.  Circumstellar dust shells form in relatively pure environments, but interstellar dust represents a mixture of material ejected by many generations of evolved stars with a wide variety of progenitor masses.  Since oxygen-rich dust shells outnumber carbon-rich dust shells, oxygen-rich dust dominates in the interstellar medium.  This means that an oxygen-rich dust spectrum can arise in either a pre-main-sequence or a post-main-sequence environment, but carbon-rich spectra will only appear in post-main-sequence objects.  \\subsubsection{Oxygen-Rich Dust Emission \\label{sec.orich}}  The oxygen-rich post-main-sequence objects can be organized along the sequence from AGB source $\\rightarrow$ OH/IR source $\\rightarrow$ PPN $\\rightarrow$ PN.  This sequence assumes that as the average oxygen-rich star evolves up the AGB, its mass-loss rate increases, eventually enshrouding it so deeply within its circumstellar dust shell that it disappears completely from the optical sky.  This first stage of development is well-documented.  \\cite{jon90}, in a study of variable AGB sources identified by the Air Force Geophysics Laboratory (AFGL) infrared sky survey \\citep{pm83}, showed that as Miras evolve to OH/IR stars, the period of variability and mass-loss rate steadily increase as the infrared colors progressively redden.  They showed photometrically that this evolution transformed the silicate emission feature at 10~\\mum\\ to a deep absorption feature.  \\cite{le90} illustrates this change spectroscopically with LRS data in his Figure 3.  Examining the composition of each subgroup (in terms of the fraction represented by AGB sources, OH/IR stars, PPNe, and PNe) helps to organize the spectral subgroups defined in our classification system into an evolutionary sequence.  In some cases, the composition of a subgroup is obvious; in others the small sample sizes and the inherent selection effects of the SWS database limit the usefulness of this method.  The initial steps are relatively straightforward to interpret. A star on the early AGB will appear as a naked star with absorption bands from CO and SiO (1.NO).  Reinterpreting Figure 8 of \\cite{sp95} in terms of the subgroups defined here, the shift from 1.NO to 2.SE occurs between (time-averaged) spectral types of M4 and M5.  This marks the onset of significant mass loss and dust formation, but the detailed evolution through the various classes of silicate emission (broad, structured, and classic) is more difficult to trace. \\cite{sp95,sp98} found few correlations between spectral type, variability class, and the shape of the silicate dust spectrum.  They suggested that the formation of multiple shells might cloud the picture, and that the shape of the spectrum might depend on photospheric C/O ratio, which would imply that dredge-ups of processed material from the stellar interior might determine the shape of the infrared dust features. Detailed analysis of the shape of the silicate feature and related features, such as the CO$_2$ lines and the 13 and 19.5 \\mum\\ bumps, in the 2.SE subgroups may shed further light on this subject \\citep{sgkp02}.  As the dust contribution grows to dominate the stellar spectrum, the spectrum will shift from Group 2 to Group 3 (3.SE).  It will then develop into a 3.SB spectrum as the optical depth of the silicate dust increases and drives the 10~\\mum\\ feature into self-absorption. The 3.SE sources are a mixture of M stars on the AGB, M supergiants, and optically enshrouded OH/IR stars.  The 3.SB sources are more evolved, with later spectral types and longer periods of variability.  All of the sources in 3.SE were initially discovered in infrared surveys, and all show OH masers.  The lack of a single source with a 2.SB spectrum suggests that the transition to self-absorption occurs in Group 3.  Evolution continues from 3.SB to 4.SA, where the silicate feature goes into full absorption and the SED becomes redder.  The sources in this group are associated with OH/IR stars, many of them Mira variables, and PPNe.  The transition to SA does not occur within Group 3, because all of the 3.SA spectra are associated with pre-main-sequence sources.  It also does not appear to occur frequently within Group 4, as most of the 4.SB sources appear to be much more evolved PNe.  Rather, the transition from SB to SA appears to coincide with the transition from Group 3 to Group 4.  The stages following 4.SA are much less clear.  Ultimately, the high mass-loss rates associated with the end of the AGB-OH/IR phase will strip the envelope from the core, producing a PPN.  As the remnant shell expands and thins, revealing the ionized central regions, the source becomes a PN. How does this process manifest itself into the subgroups not yet included in the sequence:  4.SB, 4.SE, 4.SEC, and 4.SC? All four subgroups show roughly the same percentage of clearly identified PNe (58-60\\% of the sample, excluding high-mass objects and pre-main-sequence objects), but only 4.SC and 4.SEC include any sources identified as OH/IR stars or still on the AGB. Because of this,  we suspect that 4.SC and 4.SEC precede 4.SE and 4.SB on a typical evolutionary path.  \\cite{wat96} first identified crystalline species of silicates in the spectra of circumstellar dust shells associated with evolved stars using data from the SWS.  They noted that the crystalline features do not appear until the color temperature of the shell has decreased to $\\sim$300 K.  Further study of SWS data by \\cite{syl99} relates the presence of crystalline features with the optical depth at 10~\\mum.  They suggest that crystalline silicates do not appear until the mass-loss rate has crossed a certain threshold.  Thus as mass-loss rate increases, absorption strength at 10~\\mum\\ grows stronger, and color temperature reddens, a typical source will evolve to SA and then to SC.  Most of the sources in subgroups 4.SC, 4.SEC, 4.SB, and 4.SE appear in the upper middle of the HR diagram (spectral class B, A, F, and G, usually with emission lines, luminosity class I-II).  Whatever their precise order, most or all of the post-main-sequence sources with these classes of spectra are obviously in transition from the AGB or red supergiant phase to later stages of evolution.  It is likely that the difficulty in ordering these subgroups results from the wide range of stellar masses which can produce oxygen-rich dust shells (from less than 1 M$_{\\sun}$ to beyond 50 M$_{\\sun}$).  The more massive stars do not follow the standard evolutionary scenario; instead they evolve onto the super-AGB \\citep[e.g.][]{gbi94}.  Initial masses $\\ge$ 11 M$_{\\sun}$ produce final core masses beyond the Chandrasekhar limit and become supernovae.  Masses $\\ga$ 50 M$_{\\sun}$ are associated with the luminous blue variables \\citep[e.g.][]{hd94}, some of which are in the SWS sample.  With all of these sources producing oxygen-rich dust shells, perhaps it is not a surprise that the redder spectra cannot be ordered into a smooth sequence.  Another complication is the mixture of young and old sources in Groups 3 and 4 (in contrast to the oxygen-rich spectra in Groups 1 and 2 (1.NO and 2.SE), most of which are evolved sources).  In subgroup 4.SE, 9 of 24 sources are clearly pre-main-sequence; all are Herbig Ae/Be stars except for one source classified as F0e.  This represents the majority of the young sources in the sample, but three more Herbig Ae/Be stars appear in subgroup 3.SE (out of 21), both 3.SA spectra are pre-main-sequence Be stars, subgroups 4.SEC and 4.SB each contain two pre-main-sequence sources (out of 10 and 7 respectively), and one of the 14 4.SC sources is young (a T Tauri star).  Three of the four young sources in subgroups 4.SEC and 4.SB are Herbig Ae/Be stars; the other source is an Ae star.   It is unfortunate that young and old sources appearing in the same part of the HR diagram, luminous Be, Ae, Fe, and G stars, exhibit similar infrared spectral characteristics.  Determining whether these sources were evolving {\\em to} or {\\em from} the main sequence has been a long-standing problem in astronomy. \\cite{wcv89} showed that some types of young and old stars could be separated into different zones using \\iras\\ color-color diagrams. As Figures 4 and 5 show, the SWS spectra extend sufficiently beyond 20 \\mum\\ for the shape of the continuum to be defined, thus showing the underlying dust temperature. The shape of the dust continuum from $\\sim20$ to $\\sim40$ \\mum\\ might be one way of separating the young and old objects. Potentially, the more detailed Level 3 classification will address this issue.  \\subsubsection{The Carbon-Rich Dust Sequence \\label{sec.crich}}  While oxygen-rich dust can occur in both evolved stars and environments associated with star formation, carbon-rich dust only occurs in the vicinity of carbon stars or in planetary nebulae which have presumably evolved from carbon stars.  Furthermore, the range of stellar masses which evolve to carbon stars is limited to $\\ga$ 2 M$_{\\sun}$ and less than several M$_{\\sun}$ \\citep[cf.][and references therein]{wk98}.  Perhaps for these reasons, the carbon-rich spectral classes defined here fall into a reasonably ordered evolutionary sequence:  1.NC $\\rightarrow$ 2.CE $\\rightarrow$ 3.CE $\\rightarrow$ 3.CR $\\rightarrow$ 4.CR $\\rightarrow$ 4.CN.  As the mass-loss rate from a naked star with a carbon-rich photosphere (1.NC) grows, its infrared spectrum develops a strong emission feature at $\\sim$11.5~\\mum\\ from SiC, producing a 2.CE spectrum.  Further increases in mass-loss rate lead to a cooler, optically thicker shell which enshrouds the central star.  The spectrum is then classified as 3.CE.  It next evolves to 3.CR as the emitting layer of the dust shell cools and amorphous carbon begins to dominate the spectrum.  Further thickening of the dust shell shifts the spectrum to 4.CR.  The next stage is less certain because the relation of 4.CT to the sequence is not clear.  Perhaps spectra evolve from 4.CR to 4.CN (i.e. to carbon-rich PPNe), and only some unusual circumstances lead to the development of a 4.CT spectrum.  Possibly,  all carbon-rich sources pass briefly through this stage.  However, the latter possibility seems unlikely given the difficulty of fitting the 4.CT spectra into the rest of the carbon sequence."
    },
    "0201/astro-ph0201441_arXiv.txt": {
        "abstract": "We present the results from two radio integrations at 8.4 GHz using the VLA.  One of the fields, at 13$^h$+43$^\\circ$ (SA13 field), has an rms noise level of $1.49~\\mu$Jy and is the deepest radio image yet made.  Thirty-four sources in a complete sample were detected above $7.5~\\mu$Jy and 25 are optically identified to a limit of I=25.8, using our deep HST and ground-based images.  The radio sources are usually located within $0.5''$ (typically 5 kpc) of a galaxy nucleus, and generally have a diameter less than $2.5''$.  The second field at 17$^h$+50$^\\circ$ (Hercules Field) has an rms noise of $35~\\mu$jy and contains 10 sources.  We have also analyzed a complete flux density-limited sample at 8.4 GHz of 89 sources from five deep radio surveys, including the Hubble deep and flanking fields as well as the two new fields.  Half of all the optical counterparts are with galaxies brighter than I=23 mag, but 20\\% are fainter than I=25.5 mag.  We confirm the tendency for the micro-Jansky radio sources to prefer multi-galaxy systems.  The distribution of the radio spectral index between 1.4 and 8.4 GHz peaks at $\\alpha\\approx -0.75~(S\\sim\\nu^{+\\alpha}$), with a median value of $-0.6$.  The average spectral index becomes steeper (lower values) for sources below $35~\\mu$Jy, and for sources identified with optical counterparts fainter than I=25.5 mag.  This correlation may suggest that there is an increasing contribution from star-burst galaxies compared to active galactic nuclei (AGNs) at lower radio flux densities and fainter optical counterparts.  The differential radio count between 7.5 and 1000 $\\mu$Jy has a slope of $-2.11\\pm 0.13$ and a surface density of 0.64 sources (arcmin)$^{-2}$ with flux density greater than $7.5~\\mu$Jy. ",
        "introduction": "Deep surveys of the extragalactic sky have been made at X-ray (Hasinger et al.\\ 1998; Cowie et al.\\ 2001), optical (Lilly et al.\\ 1996; Madau et al.\\ 1996; Steidel et al.\\ 1999), infra-red (Goldschmidt et al.\\ 1997; Rowan-Robinson et al.\\ 1997; Elbaz et al.\\ 1999), sub-mm (Hughes et al.\\ 1998; Barger et al.\\ 1999; Blain et al.\\ 1999; Barger et al.\\ 2000; Scott et al.\\ 2000), as well as radio wavelengths (Windhorst et al.\\ 1984; Windhorst et al.\\ 1985; Condon \\& Mitchell 1984; Donnelly et al.\\ 1987; Fomalont et al.\\ 1991; Windhorst et al.\\ 1993; Richards et al.\\ 1998; Richards et al.\\ 1999; Prandoni et al.\\ 2001).  Much of the interest in these surveys of weak objects concern the evidence for star formation and AGN evolution in galaxies in the early universe, at redshifts in the range 1 to 3 (Condon \\& Yin 1990; Rowan-Robinson et al.\\ 1993; Cram 1998; Mobasher et al.\\ 1999; Steidel et al.\\ 1999; Haarsma et al.\\ 2000).  The radio observations are unique in that at radio wavelengths it is possible to peer through the gas and dust that obscures the nuclear regions of galaxies at other wavelengths.  The radio observations also have sufficient angular resolution to distinguish between emission that is driven by star formation and that driven by AGN.  Although star forming activity in the early universe is perhaps most readily observed at sub-millimeter wavelengths, current sub-millimeter instruments do not have sufficient angular resolution to avoid confusion due to blending of nearby sources, so high resolution radio observations are needed to uniquely identify optical counterparts. Moreover, because of the large negative k-correction, dusty galaxies observed at sub-mm wavelengths all have high redshifts, whereas the radio observations are sensitive to a wide range of redshift and a mix of star-burst and AGN activity.  As the most sensitive radio telescope available for these kinds of observations, the Very Large Array (VLA) has been used for many deep surveys \\citep{con84,win84,win85}.  Deep radio source surveys have also been made with the Australian Telescope Compact Array \\citep{hop99,wil00,pra01}, and with the Westerbork Synthesis Radio Telescope \\citep{gar00}, but with less sensitivity and poorer resolution than possible with the VLA.  These observations show that below levels of about 10 mJy, the number of radio sources increases more rapidly than the number between 1 mJy and 500 mJy and are composed of a different population of sources.  In this paper we report on the results of new 8.4 GHz VLA surveys which cover two fields, each of diameter $9.2'$ to the 8\\% sensitivity level of the on-axis position.  One field, located at $13^h$+43$^\\circ$, which we will designate as the SA13 field, was one of the Hubble Space Telescope Medium Deep Survey (MDS) key projects, observed in 1992-3 \\citep{gri94, win94, win95}.  These optical observations, which were made before HST refurbishment, reach a limiting sensitivity of I=25.5 mag.  The extensive complementary VLA observation detected sources as faint as $7.5~\\mu$Jy at 8.4 GHz.  A preliminary account of the HST identifications in this field \\citep{win95}, hereinafter designated as Paper I, was based on the first half of these observations which were made at lower resolution, and included only radio sources located in the small 2.5 (arcmin)$^2$ field of view of the HST.  We have also used these same VLA observations of this field to examine the small scale fluctuations in the cosmic microwave background radiation \\citep{par97}.  This field has also been observed with the VLA at 1.4 GHz, covering a much larger area, with additional KPNO optical imaging \\citep{ric00, ric02}, and with the MERLIN/VLBI National Facility of the University of Manchester at 1.4 GHz \\citep{fom02}  A second field, located at $17^h$+50$^\\circ$ and designated at the Hercules field, was previously imaged with WFPC2 in a 48-orbit exposure, and reaches a 10-$\\sigma$ point source sensitivity of about 27.7 mag in B and V and 26.8 in I \\citep{win98}. Twelve known objects in this field, including three radio-weak AGN, have spectroscopic redshifts z near 2.4 (Pascarelle et al.\\ 1996a; Pascarelle et al.\\ 1996b; Pascarelle et al.\\ 1998; Pascarelle et al.\\ 2001).  A WFPC2 image made with a medium band redshifted Ly$\\alpha$ filter indicates the presence of a number of other compact galaxy candidates with redshifts near 2.4.  The VLA image of this field reached a limiting flux density of only $35~\\mu$Jy, but provides improved statistics for radio sources above this level.  In $\\S$2 of this paper we describe the VLA observations, calibrations and imaging for the two fields.  The radio and optical parameters for the sources are given in $\\S$3, and detailed radio/optical images with discussions are given in $\\S$4, along with discussions of the identifications.  In addition, using data from five deep 8.4-GHz surveys, we discuss the spectral properties, the source density and the optical properties of the micro-Jansky radio sources.  Further discussion is given in the final section. ",
        "conclusions": ""
    },
    "0201/astro-ph0201488_arXiv.txt": {
        "abstract": "The role of thermodynamics in the evolution of systems evolving under purely gravitational forces is not completely established. Both the infinite range and singularity in the Newtonian force law preclude the use of standard techniques. However, astronomical observations of globular clusters suggests that they may exist in distinct thermodynamic phases. Here, using dynamical simulation, we investigate a model gravitational system which exhibits a phase transition in the mean field limit. The system consists of rotating, concentric, mass shells of fixed angular momentum magnitude and shares identical equilibrium properties with a three dimensional point mass system satisfying the same condition. The mean field results show that a global entropy maximum exists for the model, and a first order phase transition takes place between \"quasi-uniform\" and ''core-halo'' states, in both the microcanonical and canonical ensembles. Here we investigate the evolution and, with time averaging, the equilibrium properties of the isolated system. Simulations were carried out in the transition region, at the critical point, and in each clearly defined thermodynamic phase, and striking differences were found in each case. We find full agreement with mean field theory when finite size scaling is accounted for. In addition, we find that (1) equilibration obeys power law behavior, (2) virialization, equilibration, and the decay of correlations in both position and time, are very slow in the transition region, suggesting the system is also spending time in the metastable phase, (3) there is strong evidence of long-lived, collective, oscillations in the supercritical region. ",
        "introduction": "In contrast with ''normal'' systems with short range interactions, the thermodynamics of self-gravitating systems is non-extensive and, because of the infinite range and singularity of the Newtonian potential, cannot be treated by standard methods. A partial remedy for these problems can be constructed by confining the system in a finite volume and either using a regularized Newtonian pair interaction potential, or considering the N-body system in the mean-field limit where it is possible to construct an analytic theory of $f(\\mathbf{r},\\mathbf{v})$ , the single particle density in position and velocity. The first mean-field formulations showed that (1) spherically symmetric density profiles represent the states of highest entropy and (2) a global entropy maximum does not exist \\cite{antonov}. It is always possible to increase the entropy by simultaneously increasing the central density and transferring mass to a diffuse ''halo'' to control the value of the energy. This phenomenon, called gravothermal catastrophe in the literature, reflects the fact that an isolated and bounded gravitational system cannot be in equilibrium in the mean-field limit. Locally stable and unstable entropy extrema can exist, however, if the system's energy is above a critical value \\cite{antonov,lyndbw}, and their stability has been investigated by several authors \\cite{padna1,katz1,katz2}.  In addition to the total energy, $E$, in the mean field limit the sum of squares of the angular momentum, $L_{2}=\\lim\\limits_{N->\\infty }1/N\\sum l_{i}^{2}$ , where $\\mathbf{\\ l}_{i}$ is the angular momentum of a system element, is also an integral of motion for a spherically symmetric gravitational system, \\cite{binney} . If, along with the energy, this constraint is also included in MFT (mean field theory) the result is an anisotropic distribution in velocity. The result was first obtained by Eddington from a different route \\cite{Eddington}. Until recently, this second integral has been ignored in investigations of thermodynamic stability in the mean field (Vlasov) limit, leaving open the possibility that a centrifugal barrier could prevent collapse and stabilize the system. In fact, we have recently shown that even its inclusion in MFT cannot resolve the gravothermal catastrophe, and it persists in both the generalized microcanonical ($E-L_{2} $) and canonical ($\\beta -\\gamma $) ensembles \\cite{klinkmillL2}. The extremal solutions have the form $f\\propto \\exp (-\\beta \\epsilon )\\exp (-\\gamma l^{2})$, which coincides with the well known anisotropic density fit models that have been applied to globular cluster observations with some success(e.g. King-Michie models) \\cite {gunn,Heggie}.  In their seminal work on the gravothermal catastrophe, Lynden-Bell and Wood \\cite{lyndbw} pointed out that in the absence of the short range singularity, the possibility existed for a gravitational system to exist in different thermodynamics phases. Using mean-field models with a regularized Newtonian potential which has been dissected to remove the singularity \\cite {kiessling1} or, equivalently, introducing repulsion at short-range by imposing a local equation of state in the mean field picture \\cite {aronson,stahlkiessling2,padna1}, it has been shown that a first order phase transition can occur in both the microcanonical (MCE) and canonical (CE) ensembles. Unfortunately, for finite $N$ systems, there are no exact microcanonical results available which allow the rigorous proof of a phase transition or catastrophe. However, in the canonical case, it has been rigorously proved that the system of gravitating point masses is in a collapsed state in equilibrium in the absence of regularization \\cite {kiessling1}. Moreover, Monte-Carlo simulations for a regularized Newtonian potential confirm the gravitational phase transition in CE \\cite{sanchez} in large $N$ systems.  While mean field theories support the existence of phase transitions in gravitational systems, it is important to point out that there is no guarantee that these equilibrium states will be realized by dynamical evolution. In fact, there is no proof that the two operations of taking (1) the mean field limit, or (2) the infinite time average, commute \\cite{latora}% . Rather, simulations of the one dimensional self-gravitating system consisting of parallel mass sheets provide strong evidence to the contrary \\cite{mineufeix}. As a consequence, the relation between maximum entropy solutions of the stationary mean field equation and the time average distribution functions resulting from N body simulation, or dynamical evolution in nature, has not been fully established. This is a deep question which will not be explored further here. Thus, although much is known concerning the ''equilibrium'' properties in the mean field limit, the dynamical properties of gravitational phase transitions are not well known due to a lack of true N-body simulations, which are also important for explaining the evolution of stellar clusters, galaxies, etc.. At the present time the mean-field predictions of the gravitational phase transition have only been dynamically confirmed (in both MCE and CE) for the model system consisting of irrotational, concentric, mass shells \\cite{millyounglett}. In that model, the Newtonian singularity was screened by the introduction of an inner barrier which excluded mass from the system center.  The aim of this paper is to investigate and understand the dynamical features of gravitational phase transitions in N-body simulations for the model of a purely Newtonian system in which $l_{i}^{2}=l^{2}=L_{2}$ for each system element. We will refer to this system as the $l^{2}$ model. We will explicitly investigate a system of rotating, concentric, mass shells. In the mean field limit, this system shares important features, e.g. the equilibrium density and radial velocity distribution, with the more realistic system of point masses. Recently, with I. Prokhorenkov, we showed that the gravitational phase transition is present in both MCE and CE in certain regions of $l^{2}$ and we studied its properties using mean field theory \\cite{klinkmillprok}. Moreover, we rigorously proved the existence of an upper bound for the entropy in the MCE (lower bound in the free energy in the CE) for $l^{2}\\neq 0$ in the same work. It is important to understand that this model demonstrates the significance of the influence of angular momentum on the thermodynamics of self-gravitating systems, even if they are spherical. The generalized microcanonical ($E-L_{2}$) and canonical ($\\beta -\\gamma $) ensembles discussed earlier \\cite{klinkmillL2}are appropriate for large $N$ spherical systems where angular momentum exchange occurs. While providing the most general mean field description \\cite{klinkmillL2}, $L_{2}$ is still not a sufficient constraint to resolve the gravothermal catastrophe. Clearly the culprit is angular momentum exchange, which is still permitted even if both energy and $L_{2}$ are fixed, and allows the transfer of mass to the system center. An open question is the possible existence of additional mechanisms for establishing a centrifugal barrier which prevents, or strongly inhibits, collapse in nature, e.g. in globular clusters or molecular clouds. This will be taken up in the final section.  In the following, first we briefly review the $l^{2}$ model, including the mean-field predictions, and discuss the main features of the N-body algorithm we designed for its dynamical simulation. Next we turn to the simulation results in a region of the $(E,l^{2})$ phase plane containing the microcanonical phase transition region. In each phase we compare the time-averaged equilibrium properties with the predictions of mean field theory resulting from our earlier investigation, and also touch on finite size effects. We then go on to study both equilibrium and dynamical features which cannot be predicted by mean field theory, such as the variance of fluctuations, and correlations in both time and position. In addition to the system behavior near the phase transition, we pay particular attention to the critical point, and the supercritical region. Finally we consider the surprising features exhibited by different stages of the relaxation process itself, and their dependence on energy, $l^{2}$, and population, and discuss the possible presence of collective modes.  \\begin{figure}[tbp] \\centerline{\\includegraphics[scale=0.5]{FIG01A}} \\caption{Plot of the entropy extrema in the microcanonical ensemble for $l^{2}=5\\times 10^{-5}$ in the mean-field limit. As we can see, in the transition region multiple solutions are present, and we marked the solutions for $E=2$ with labels 1, 2, and 3. Only phase 1 (quasi-uniform phase) is a global entropy maximum above the transition point ($E=1.9)$, while phase 2 (condensed phase) is locally stable, and phase 3 is a saddle point. } \\label{figtransmean} \\end{figure}  \\begin{figure}[tbp] \\centerline{\\includegraphics[scale=0.5]{FIG01B}} \\caption{Density profiles of the three distinct phases labelled in Fig.~\\ref {figtransmean} for $E=2$ in the transition region. Above the transition point only the quasi-uniform phase (phase 1) is globally stable, while the condensed phase (phase 2) is locally stable. The third entropy extremum (phase 3) is unstable.} \\label{figdensityprof} \\end{figure}  \\begin{figure}[tbp] \\centerline{\\includegraphics[scale=0.5]{FIG02}} \\caption{The mean-field microcanonical phase diagram. The system can exist in 3 different types of phase depending on the energy and $l^{2}$, and we named these phases in analogy with normal systems. We also indicated the metastable regions which are important for understanding the dynamical behavior of the system near the transition point.} \\label{figphasediagram} \\end{figure}  \\begin{figure}[tbp] \\centerline{\\includegraphics[scale=0.5]{FIG03}} \\caption{The time averaged relative populations $\\bar{P}_{i}$ of relaxed systems in the high energy region at $E=4$ and $l^{2}=5\\times 10^{-5}$ with $% N_{bin}=20$. The bins were obtained from the equal-mass radii of the mean-field equilibrium density profile. We observe good convergence to the mean-field density profile with increasing $N$. As a result of finite size effects, the central density is higher than the mean-field density profile in the high energy region, and causes a reduction in the virial ratio.} \\label{figdensityhigh} \\end{figure}  \\begin{figure}[tbp] \\centerline{\\includegraphics[scale=0.5]{FIG04}} \\caption{The averaged virial ratio $\\overline{VR}$ for $56$ simulations for $% N=16, 32, 64, 128$ in and around the mean-field transition region. We also included the mean-field results for the globally stable states which show that the system undergoes a first order phase transition at $E=1.9$, where the virial ratio becomes discontinuous. The simulation results converge to the mean-field predictions with increasing $N$, but the transition point is shifted and the transition region is rounded, in good agreement with finite-size scaling.} \\label{figphasesim} \\end{figure}  \\begin{figure}[tbp] \\centerline{\\includegraphics[scale=0.5]{FIG05}} \\caption{The time correlation function of the kinetic energy $C_{KE}\\left( \\protect\\tau \\right) $ in 5 different phase regions for $N=64$, where $% \\protect\\tau $ is in units of $t_{dyn}$. In both the high energy ($% E=4,l^{2}=5\\times 10^{-5} $) and supercritical phase ($E=4,l^{2}=5\\times 10^{-3}$), correlations are smaller, while in the low energy phase ($% E=1.5,l^{2}=5\\times 10^{-5}$), at the transition point ($E=1.9,l^{2}=5\\times 10^{-5}$), and at the critical point ($E=1.052,l^{2}=1.1\\times 10^{-4}$) relatively stronger correlations can be observed. Note that only the initial decay of the correlation functions follows an exponential behavior, then a long tail develops.} \\label{figKEcorr} \\end{figure}  \\begin{figure}[tbp] \\centerline{\\includegraphics[scale=0.5]{FIG06}} \\caption{Position correlation matrix in relative bin populations $C_{ij}$, for $N_{bin}=20$ and $i>j$ in the low energy phase ($E=1,l^{2}=5\\times 10^{-5},N=64$). The bin radii represent the equal-mass layers obtained from the mean-field density profile. Strong anti-correlation is present close to the high density central core region.} \\label{figpopcorrlow} \\end{figure}  \\begin{figure}[tbp] \\centerline{\\includegraphics[scale=0.5]{FIG07}} \\caption{Position correlation matrix in relative bin populations $C_{ij}$, for $N_{bin}=20$ and $i>j$ in the phase transition region ($% E=2,l^{2}=5\\times 10^{-5},N=64$). The bin radii represent the equal-mass shells obtained from the mean-field density profile of phase 1 in Fig.~\\ref {figtransmean}. Strong anti-correlation is present between the center and the remainder of the system.} \\label{figpopcorrphase} \\end{figure}  \\begin{figure}[tbp] \\centerline{\\includegraphics[scale=0.5]{FIG08}} \\caption{Position correlation matrix in relative populations $C_{ij}$, for $% N_{bin}=20$ and $i>j$ near the critical point of the phase diagram ($% E=1.052,l^{2}=1.1\\times 10^{-4},N=64$). The bin radii represent the equal-mass shells obtained from the mean-field density profile. The behavior is similar to that of the low energy phase, although the anti-correlating region is wider. This reflects the fact that, at the critical point, the difference between the two distinct phases vanishes. This intermediate state still has a core-halo structure, but with a more extended core region.} \\label{figpopcorrcrit} \\end{figure}  \\begin{figure}[tbp] \\centerline{\\includegraphics[scale=0.5]{FIG09}} \\caption{Violent relaxation in the high energy phase ($E=4,l^{2}=5\\times 10^{-5}$) in four simulations with $N=16, 32, 64 , 128$. In the figure $N$ is increasing from bottom to top. We observe that virialization is faster with increasing $N$. For $N=64, 128$ the differences become marginal as we approach the mean-field limit.} \\label{figvirrelaxhigh} \\end{figure}  \\begin{figure}[tbp] \\centerline{\\includegraphics[scale=0.5]{FIG10}} \\caption{Violent relaxation in the low energy phase ($E=1,l^{2}=5\\times 10^{-5}$) at $N=16, 32, 64, 128$ (from top to bottom in the figure). Here the virialization is also fast for larger values of $N$, and takes place on a much shorter time-scale than in the high energy phase. Note that for $% N=64, 128$, the system virializes in $100t_{dyn}$. } \\label{figvirrelaxlow} \\end{figure}  \\begin{figure}[tbp] \\centerline{\\includegraphics[scale=0.5]{FIG11}} \\caption{Violent relaxation in the phase transition region ($% E=2,l^{2}=5\\times 10^{-5} $) in four different simulations. From bottom to top $N=16, 32, 64, 128$. Compared with the high and low energy phases, the relaxation is slower for larger $N$ as opposed to the process in the high and low energy phases. As we approach the mean-field limit, the phase transition becomes sharper and the increasing influence of the meatstable phase slows down the relaxation.} \\label{figvirrelaxphase} \\end{figure}  \\begin{figure}[tbp] \\centerline{\\includegraphics[scale=0.5]{FIG12}} \\caption{Plots of $\\protect\\sigma _{r}^{2}$ vs. time which indicate the relaxation process in $\\protect\\mu $ space for simulations with $N=16, 32, 64, 128$ in the high energy phase ($E=4,l^{2}=5\\times 10^{-5}$). In order to avoid statistical error while sampling the populations, we chose snapshot times $t_{s}=0.125, 0.25, 0.5, 1$ for the corresponding values of $N$. For all $N$, the same power-law type of relaxation, $\\protect\\sigma % _{r}^{2}\\propto 1/t$, can be observed.} \\label{figrelaxhigh} \\end{figure}  \\begin{figure}[tbp] \\centerline{\\includegraphics[scale=0.5]{FIG13}} \\caption{Plots of $\\protect\\sigma _{r}^{2}$ vs. time which indicate the relaxation process in $\\protect\\mu $ space with simulations $N=16, 32, 64, 128$ in the phase transition region ($E=2,l^{2}=5\\times 10^{-5}$). In order to avoid statistical error during sampling the populations, we chose snapshot times $t_{s}=0.125, 0.25, 0.5, 1$ for the corresponding $N$. The power-law behavior can still be observed, but large fluctuations start to take place with increasing $N$. As a result of the influence of the metastable phase, relaxation is slower for larger $N$ in this region.} \\label{figrelaxphase} \\end{figure}   \\begin{figure}[tbp] \\centerline{\\includegraphics[scale=0.5]{FIG14}} \\caption{Plots of $\\protect\\sigma _{r}^{2}$ vs. time which indicate the process of relaxation in $\\protect\\mu $ space for simulations with $N=16, 32, 64, 128$ in the low energy phase ($E=1,l^{2}=5\\times 10^{-5}$). In order to avoid statistical error during sampling the populations, we chose snapshot times $t_{s}=0.125, 0.25, 0.5, 1$ for the corresponding values of $% N $'s. In the high energy phase, we observe a similar power-law relaxation which also follows $\\protect\\sigma _{r}^{2}\\propto 1/t$.} \\label{figrelaxlow} \\end{figure}  \\begin{figure}[tbp] \\centerline{\\includegraphics[scale=0.5]{FIG15}} \\caption{The time correlation of the kinetic energy $C_{KE}\\left( \\protect% \\tau \\right) $ versus time for three regions of the supercritical phase at the energy $E=0$. With increasing $l^{2}$, an oscillating part starts to sit on top of the correlation functions which suppresses the early exponential behavior, and a very long tail starts to develop.} \\label{figKEcorrfluidl2} \\end{figure}  \\begin{figure}[tbp] \\centerline{\\includegraphics[scale=0.5]{FIG16}} \\caption{$N$ dependence of the oscillations in $C_{KE}\\left( \\protect\\tau % \\right) $ at $E=0$ and $l^{2}=0.1$. As we approach the mean-field limit, oscillations persist for a longer time and correlations in the tail grow.} \\label{figKEcorrfluidN} \\end{figure} ",
        "conclusions": "The role of thermodynamics in controlling the evolution of gravitational systems is only partially understood, and there are many open questions. It is our impression that recently this subject is attracting increased attention. Observations of the radial density dependence of globular clusters show that they fall into two groups, with either a smoothly decreasing density profile with increasing radius, $r$, or a sharp central peak and a more diffuse halo \\cite{Heggie}. The evidence suggests that a thermodynamic interpretation may be possible, i.e. perhaps the clusters can exist in different phases at observable times \\cite{Heggie}.  Here we investigated the dynamics of the gravitational phase transition in the $l^{2}$ model of a self-gravitating system in the MCE in which the mass elements are thin, concentric, shells. In this idealized model, instead of regularizing the singularity of the Newtonian potential, we simply fixed the square of the angular momentum of each shell. In section \\ref{Nbody}, we showed that, in the mean field limit, the equilibrium states of this system correspond to those of the more realistic system of point masses when the same constraint is imposed. In an earlier work, we showed that a correspondence also exists under less restrictive circumstances \\cite {klinkmillL2}. The constraint of constant $l^{2}$ for each shell, or particle, establishes a centrifugal barrier which prevents the occurrence of the gravothermal catastrophe, and induces a first order phase transition in both CE and MCE \\cite{klinkmillprok}.  One important conclusion which can be reached from this study is that angular momentum exchange plays an important role in a self-gravitating system which alters the thermodynamic behavior. Another is that it is not necessary to soften the singularity in the Newtonian potential to obtain a transition. The effect of the singularity can be blocked by other mechanisms. This is especially relevant in stellar clusters, where the distance between stars is too great for softening to be important. If stellar systems can exist, or even approximate, different thermodynamic phases, the influence of the singularity needs to blocked, at least temporarily. In our earlier, mean field, study of a spherical system which permitted exchange, no transition was found and the gravothermal catastrophe could not be prevented in spite of the fact that both $L_{2}$ and $E$ were fixed. At this time it is not clear how this might occur. One possibility is that angular momentum exchange between stars may occur so slowly that, in the present epoch, thermodynamics could be influenced by an effective centrifugal barrier. Since globular clusters are approximately spherical, this is worth investigating. Another possibility is that stellar interactions with hard binaries in the cluster core may also establish a centrifugal barrier. These conjectures require further investigation.  In this work we used N-body dynamical simulation to verify the earlier mean-field predictions \\cite{klinkmillprok}. Our N-body simulations confirmed the mean-field phase transition, which was shifted and broadened due to finite size effects. We also verified that finite size scaling is in very good agreement with the observed shifting and rounding in the transition region. An interesting feature of the convergence to MFT was the observation that agreement in the density with increasing $N$ was much slower near the system center for small values of $l^{2}$. This appears to be a discreteness effect. For a small value of $l^{2}$ the density is changing rapidly due to the competition between the centrifugal barrier and the largely unscreened gravitational potential, which combine to form a very narrow minimum in the effective potential. Because of discreteness effects, the mean local density cannot change this rapidly in the dynamical simulation, and this was readily observable in the average population of the central bin.  In addition to confirming agreement with mean field theory for the time average of physical quantities, such as the density and the kinetic and potential energies, we also studied fluctuations in density and kinetic energy, correlations in position, and the dynamical behavior of the system in each phase, the transition region, and at the critical point, i.e. properties which are not accessible from MFT. These included virialization, relaxation to equilibrium, and the decay of fluctuations in kinetic energy through its autocorrelation function. In general, as $N$ becomes large, it is expected that the Vlasov regime will be approached. In this regime $% f(r,v) $ completely characterizes the system at each time. It was shown long ago that in a system with a smooth, bounded, potential, fluctuations and correlations decay in the Vlasov limit \\cite{braunhepp}. In the large $N$ scaling regime it is easy to show that the variance of fluctuations in both bin populations and total kinetic energy should decay as $N^{-1}$. We found this behavior over the complete range of population for the former, while a wide variation of power law exponents characterized the reduction of the kinetic energy variance depending on the thermodynamic state. This suggests that our simulations were not fully in the scaling regime for the smaller populations considered. This is not surprising, and was supported by the size of the spontaneous kinetic energy fluctuations as well as the $N$ dependence of the remaining dynamical quantities.  In the transition region, strong corroborating evidence was obtained to support the ansatz that the system is fluctuating in time between the stable and metastable phase. First of all, relaxation to equilibrium as measured by the $\\sigma _{r}^{2}$ statistic takes longer and, at a given time, the fluctuations are much larger than those occurring away from the transition. Second, virialization occurs on the same time scale as $\\sigma _{r}^{2}$, whereas at low energy, the virialization takes place in about $100t_{dyn}$. Fast virialization is also observed in other model systems which lack a transition, such as the planar sheet system\\cite{hohlbroadus,toshio2}, and point mass systems \\cite{binney, Heggie} where it occurs in 50-100 dynamical times, long before relaxation to equilibrium has occurred. Third, near the transition an extremely long time tail appears in the kinetic energy autocorrelation function, indicating that the system continues to feel the presence of the metastable phase. The lack of oscillation in $C_{KE}(t)$ suggests that a slow, diffusive, mode is dominating the linear relaxation.  The observation that relaxation to equilibrium exhibits power law behavior was unanticipated and is crucial for understanding the system dynamics. In the Astrophysics literature it is usual to define a relaxation time for gravitational evolution \\cite{Heggie,binney} as the time it takes for a typical star to be sufficiently deflected that the change in its velocity is on the order of its mean. The standard result is $t_{relax}\\approx (N/10\\ln N)t_{dyn}$. However, here we see that relaxation does not occur with a characteristic time, but rather is scale free, with $\\sigma _{r}^{2}\\sim t^{-1}$. An open question is whether relaxation is scale free in higher dimension as well, e.g. for a point mass system.  There are now two gravitational systems in which the existence of a phase transition has been demonstrated both by mean field theory and, to within finite scaling, dynamical simulation. The system of irrotational shells studied earlier \\cite{millyounglett} and the system considered here share many similarities, but differ in one feature. Each system is characterized by an additional parameter besides the energy: In the non-rotating model, the inner barrier radius $a$ plays the role of $l^{2}$ so its maximum density occurs at $r=a$. Each of these systems exhibits a transition in MCE to a more centrally condensed phase as the energy is lowered. Each system has a critical value of the new parameter above which the transition is not allowed. However, in the present system we find evidence of persistent fluctuations that develop long lived oscillations as $N$ is increased in MCE, and local correlation in position near the system center. This suggests that collective oscillations may be occurring in the system. Persistent oscillations were also observed in simulations of the planar sheet system for large $N$ \\cite{rouet}.  We are currently modelling the system via dynamical simulation for the canonical ensemble. This is accomplished by introducing thermalizing collisions at the outer boundary. There are strong contrasts with the MCE which we will report in a separate work. We also plan to use dynamical simulation to investigate the locally stable regime which arises when angular momentum exchange is allowed \\cite{klinkmillL2}. In addition, we would like to suggest that Vlasov dynamics may give insight into the collective behavior found in the fluid phase of the $l^{2}$ model."
    },
    "0201/astro-ph0201394_arXiv.txt": {
        "abstract": "The Princeton Variability Survey (PVS) is a robotic survey which makes use of readily available, ``off-the-shelf'' type hardware products, in conjunction with a powerful set of commercial software products, in order to monitor and discover variable objects in the night sky.  The main goal of the PVS has been to devise an automated telescope and data reduction system, requiring only moderate technical and financial resources to assemble, which may be easily replicated by the dedicated amateur, a student group, or a professional and used to study and discover a variety of variable objects, such as stars.  This paper describes the hardware and software components of the PVS device as well as observational results from the initial season of the PVS, including the discovery of a new bright variable star. ",
        "introduction": "Amateur astronomers have long been involved in the study of variable stars.  Organizations like the AAVSO (http://www.aavso.org) and the Center for Backyard Astrophysics (http://www.astro.bio2.edu/cba) have promoted the study of a variety of phenomena by amateurs and overseen the collection of many years of high quality data.  As technology, particularly the CCD, has become accessible to a wider ranger of enthusiasts, the opportunities for the amateur astronomer to conduct or lead scientific inquiry have grown tremendously.  Recent projects such as TASS (http://stupendous.rit.edu/tass/tass.shtml) and  HCRT (http://www.mtco.com/\\~~jgunn) have begun to make significant contributions along these lines. The importance of a comprehensive knowledge of variability across the entire sky has been outlined by Paczy\\'nski (2001) and explored experimentally through projects such as ASAS (Pojmanski 1997, 1998).  Approximately $90\\%$ of all bright, existing variable stars are thought to be yet undiscovered, with ASAS finding that approximately $3\\%$ of a set of 140,000 monitored stars turned out to be newly discovered variables with a wide range of periodicities, as discussed by Eyer \\& Blake (2001).  The potential contribution from the study of a set of several thousand stars, even for just a few nights, is enormous.  The ultimate goal of the PVS has been to design a system to be utilized for just this purpose.  In order for the desired scientific results to be achieved in as simple a manner as possible, a number of criteria must be met.  Availability of components, ease of implementation, and reproducibility of results were all key concerns during the development of the PVS system.  Today, a wide array of advanced devices are available to the savvy amateur and when these devices are combined with powerful and inexpensive computational resources, the ability to survey the sky from a backyard is shown to be well within reach.  The optical and electronic components of the PVS are easily available through mail order, or even a local retailer, and products already widely owned by amateurs were integrated whenever possible. Though the commercial hardware systems utilized in the PVS are shown to perform well, the power of the PVS design is derived from the software that controls all aspects of data collection and actively compensates for mechanical and electronic shortcomings in the hardware components.  This collection of commercially available software products, all adhering to a set of inter-program cooperation guidelines called ASCOM (http://ascom-standards.org), work together on a PC, under Microsoft Windows, to control a telescope, camera, and, if required, a telescope dome.  These programs allow simple procedures to be controlled from a Graphical User Interface (GUI) and for complex observing sessions and data reductions to be scripted in Perl, Visual Basic Script, or C++ languages.  This paper provides descriptions of the hardware and software utilized in the PVS so that a similar system may be constructed by interested parties. Examples are presented of the type and quality of results that may be expected. ",
        "conclusions": "The design of a robotic photometric telescope, built out of commercially available and inexpensive components, is outlined. Results from one month of observations with the PVS instrument are presented, and it is demonstrated that the instrument produces high quality photometric results, even from a location with extremely bad atmospheric conditions.  The discovery of a new short period variable star is presented.  Though the system performed as well as could be expected given the atmospheric conditions at the Princeton, New Jersey, site, the full potential of the PVS design cannot be determined until it is tested at a site with better conditions.  A decrease in sky brightness by 1 magnitude, an improvement of mean seeing to 3'', and a higher percentage of clear nights would result in a tremendous increase in the number of variable stars that could be discovered by the PVS system.  Many suitable sites exists, including ones, such as Mt.Wilson, within easy reach of major metropolitan areas of the American South West. Instrumental and computational techniques are developed and tested here so that a system  could be efficiently and economically installed and operated at a better site."
    },
    "0201/astro-ph0201111_arXiv.txt": {
        "abstract": "{We have detected 18 sources over 6 $\\sigma$ threshold within two regions 8$\\arcmin$.3$\\times$16$\\arcmin$.9 and 8$\\arcmin$.3$\\times33\\arcmin$.6 in the vicinity of the point with $\\alpha$=03$^h$31$^m$02$^s$.45 (J2000) and $\\delta$=+43$\\degr$47$\\arcmin$58$\\arcsec$.5 (J2000) using a $CHANDRA$ ACIS (S+I) observation. Two of the sources were detected before with $ROSAT$ HRI and one source could be closely identified with a star in the optical catalog, USNO A-2. We have also studied source spectra applying four spectral models to the data. Most of the sources can be classified as Cataclysmic Variable, Low Mass X-ray Binary or single star candidates due to their spectral characteristics and luminosities. We also searched for the extragalactic origin for these 18 sources. The source count rates vary between 5.8$\\times 10^{-4}$-4.7$\\times 10^{-3}$ counts/s. Due to low count rates temporal characteristics of the sources can not be studied effectively. ",
        "introduction": "The imaging capability and high sensitivity of the $Chandra$ $X-Ray$ $Observatory$ offers significant advantages for source detection. The good instrumental response for photon energies up to 8 keV allows detection of faint sources even if the source is subject to high absorption. The main goal of this search was to detect faint and Super Soft X-ray Sources (SSS) using a 95 ks observation obtained by the $CHANDRA$ Advanced CCD Imaging Spectrometer (ACIS) detector. The original target of the observation was a classical nova Persei 1901 (GK Per; principal investigator= \\c{S}. Balman). The scientific aim of the proposal was deriving and studying the spectrum of the first classical nova shell resolved and detected in the X-ray wavelengths (Balman \\& \\\"Ogelman \\cite{bogel}). The $CHANDRA$ data reveals for the first time a classical nova evolving like a young, miniature supernova remnant. The scientific results of the $CHANDRA$ observation can be found in Balman (\\cite{balman1}, \\cite{balman2}). \\\\ As expected most of the sources we detected were faint sources, but we did not find any SSS above 5$\\sigma$ confidence level. We analyzed the spectra of these new X-ray sources detected in two regions 8$\\arcmin$.3$\\times$16$\\arcmin$.9 and 8$\\arcmin$.3$\\times33\\arcmin$.6 around the vicinity of the point with $\\alpha$=03$^h$31$^m$02$^s$.45 (J2000) and $\\delta$=+43$\\degr$47$\\arcmin$58$\\arcsec$.54 (J2000) within an energy range of 0.3-10 keV in order to identify their nature.\\\\ Details of the observation and the analysis methods are given in Section 2. Section 3 discusses the detection properties of the sources and Section 4 is on the spectral analysis and discussion.\\\\ ",
        "conclusions": "The source spectra and insturment responses are generated using CIAO (version 2.0) and analyzed using XSPEC (version 11.0.1) applying four models to the source data between 0.3 and 10 keV: blackbody, bremsstrahlung, powerlaw and VMEKAL with absorption. Spectra were grouped to have 5-10 counts per energy bin. The absorbed fluxes are found to be in a range 1$\\times$10$^{-16}$ - 3$\\times10^{-14}$ erg cm$^{-2}$ s$^{-1}$.\\\\ Table-2 shows the spectral parameters of the blackbody and power law models and Table-3 shows the spectral parameters of the bremsstrahlung and VMEKAL models. The errors on the spectral parameters are at 2$\\sigma$ confidence level. Figure 1 shows the fitted source spectra of all the 18 sources that has the best $\\chi^2_\\nu$s. The residuals in the figures are omitted since they were in a range of 2$\\sigma$ - {-2}$\\sigma$, and did not show significant fluctuations.\\\\ Three of the sources (Src 4,12,16) are fitted using a power law model with photon indicies between 2.5 - 4 and have blackbody temperatures around 0.5 keV resembling the spectra of the Anomolous X-Ray Pulsars (AXPs) (Israel et al.\\cite{israel}). At a distance of 10 kpc, the luminosities of these three sources are found to be around 10$^{32}$ ergs/s, and thus we exclude the possible AXP connection because the X-ray luminosities of AXPs are around 10$^{35}$ ergs/s. Assuming a source distance of 10 kpc the luminosities of all the 18 sources are calculated to be around 10$^{32}$ - 10$^{33}$ ergs/s. Such luminosities are consistent with galactic quiescent CV and quiescent LMXB origin. Two of the sources (Src 11,13) show evidence for line emmision with high absorbtion ( $N_H >$ 3$\\times$ 10$^{21}$ cm$^{-2}$) and the best fit is the absorbed VMEKAL model. These sources are strong Cataclysmic Variable and quiescent LMXB candidates (Warner \\cite{warner}; Verbunt et al. \\cite{verbunt}; Guseinov et al. \\cite{guseinov}). In addition, luminosities around 10$^{32}$ ergs/s could be attained by type O, B or giant stars as a consequence of shocks in the stellar winds or coronal emission (Cassinelli et al. \\cite{cassinelli}; Schmitt et al. \\cite{schmitt}). When we exclude the sources with bremsstrahlung temperatures above 1 keV (since almost all of O, B type or giant stars have temperatures below 1 keV), we are left with 3 candidates for galactic stars; Src 6, 9 and 12. In general, we reject an HMXB origin for our sources, since their luminosities are low compared with an HMXB where the luminosity is $\\sim$10$^{36}$ ergs/s (Guseinov et al. \\cite{guseinov}). The luminosities of the sources are around 10$^{37}$ ergs/s at 4 Mpc in agreement with the luminosities of the X-ray Binaries in other local galaxies, however we do not observe any host galaxy at those directions (Bauer et al. \\cite{bauer}). As noted in Section 1, none of these sources can be SSSs. The luminosities of the SSSs are around Eddington Luminosity (10$^{38}$ergs/s) and the blackbody temperatures are between 10-60 eV with almost no emmision above 1 keV (Kahabka \\& van den Heuvel \\cite{kahabka}). None of our 18 sources have such spectral characteristics. For the case of Dim Thermal Neutron stars (DTNs) or cooling neutron stars, the spectra of DTNs and cooling neutron stars are very soft and almost all of them are within 100 pc distance (Alpar \\cite{alpar}; \\\"Ogelman \\cite{ogelman}). A neutron star with a surface temperature around 10$^6$ K and a 10 km radius has an expected luminosity around 10$^{33}$ ergs/s. The observations on cooling neutron stars depend on the sensitivity of the X-ray telescopes. Among our 18 sources, the softest two (Src 6,11) have blackbody temperatures of 0.19 keV. This temperature is relatively high for a DTN or a cooling neutron star, however we do not exclude the possible DTN or cooling neutron star connection since the sensitivity of $CHANDRA$ allows us to observe fluxes around 10$^{-15}$ ergs cm$^{-2}$ s$^{-1}$ consistent with observation of a neutron star at 10 kpc with a temperature $\\sim$0.1 keV. Due to the relative hardness of the spectra of these 18 sources, we can also say that none of these sources are isolated hot white dwarf candidates since the temperatures of the hot white dwarfs are ultrasoft; a few eVs (Vennes \\cite{vennes}; Finley et al. \\cite{finley}).\\\\ We searched for the extragalactic origins of all the 18 sources. The rest frame luminosities of AGNs and galaxies are known to be between 10$^{39}$ and 10$^{45}$ ergs/s, and their spectra are best fitted with power law models having photon indicies around 1.7 - 2 (Brandt et al. \\cite{brandt}; Ishisaki et al. \\cite{ishisaki}). The rest frame luminosities of nearly all of the 18 sources are calculated to be around 10$^{42}$ - 10$^{44}$ ergs/s after fitting zmodels using XSPEC. Thus, we can not exclude the possibility of extragalactic origin for any of the sources. Two highly absorbed sources Src 3 and 5 have rest frame luminosities around 10$^{42}$ ergs/s (redshift=0.206) and 10$^{43}$ ergs/s (redshift=0.500) respectively. The photon indicies of these two sources are 1.61 and 1.89 respectively. Hence, we might categorize these two sources as strong AGN and galaxy candidates. Additionally, the N$_H$ values may be further evidence to exclude the extragalactic connection. The sources with N$_H$ higher than the value of the galactic N$_H$ in the direction of GK Per are more likely to be of extragalactic origin than the sources having the same order of N$_H$ with the galactic value. We may say that none of these sources are clusters since they are not extended and the rest frame luminosities of the clusters are usually higher than these values (Schindler \\cite{schindler}).\\\\ In addition, we searched for temporal characteristics of the new sources. We performed power spectrum analysis on three of the sources with the highest count rates (Src 4, 6, 18). However we could not find any significant periods. The 3$\\sigma$ - 4$\\sigma$ detection threshold of power is around 40 in our data, and the power upper limit of the three sources is found to be 20. \\\\ Drawing conclusions about the classification of all the 18 sources is difficult using the X-ray data at hand. We can not study their temporal characteristics due to low count-rates, and we do not have detailed spectra. We have planned further deep observations of this field in the optical wavelengths with the 1.5 m telescope of the National Observatory at Antalya, Turkey to ensure optical identification. Complementary observations in other wavelengths are necessary for proper classification of these 18 sources."
    },
    "0201/astro-ph0201327_arXiv.txt": {
        "abstract": "We present the critical conditions for hot trans-fast magnetohydrodynamical (MHD) flows in a stationary and axisymmetric black-hole magnetosphere. To accrete onto the black hole, the MHD flow injected from a plasma source with low velocity must pass through the fast magnetosonic point after passing through the ``inner'' or ``outer'' Alfv\\'en point. We find that a trans-fast MHD accretion solution related to the inner Alfv\\'en point is invalid when the hydrodynamical effects on the MHD flow dominate at the magnetosonic point, while the other accretion solution related to the outer Alfv\\'en point is invalid when the total angular momentum of the MHD flow is seriously large. When both regimes of the accretion solutions are valid in the black hole magnetosphere, we can expect the transition between the two regimes. The variety of these solutions would be important in many highly energetic astrophysical situations. ",
        "introduction": "In order to explain the activity of active galactic nuclei (AGNs) and compact X-ray sources, we consider a black hole magnetosphere in the center of these objects. The magnetosphere is composed of a central black hole with surrounding plasmas and a large scale magnetic field. The magnetic field is originated from an accretion disk rotating around the black hole. The electrodynamics of the black hole magnetosphere has been discussed by many authors; force-free magnetospheres were discussed in \\citet{Membrane86} and more general magnetospheres in \\citet{Punsly01}.   In the black hole magnetosphere, because of the strong gravity of the black hole and the rapid rotation of the magnetic field, both an ingoing plasma flow (accretion) and an accelerated outgoing plasma (wind/jet) should be created. The plasma would be provided from the disk surface and its corona. When the plasma density in the magnetosphere is somewhat large, the plasma inertia effects should be important. In this case, the plasma would be nearly neutral and should be treated by the ideal magnetohydrodynamic (MHD) approximation \\citep{Phinney83}, so the plasma streams along a magnetic field line, where the magnetic field line could extend from the disk surface to the event horizon or a far distant region (\\cite{Nitta-TT91}; see also \\cite{TT2001}). The outgoing flow effectively carries the angular momentum from the plasma source, and then the accretion would continue to be stationary, releasing its gravitational energy. The magnetic field lines connecting the black hole with the disk, which are mainly generated by the disk current, may not connect directly to the distant region, but via the disk's interior the energy and angular momentum of the black hole can be carried to the distant region; the energy and angular momentum transport inside the disk is not discussed here.   If the plasma density is sufficiently low and the magnetosphere is magnetically dominated, one can expect the pair-production region along open magnetic flux tubes, which connect the black hole to far distant regions directly \\citep{beskin97, Punsly01}. Further, we also expect the Blandford-Znajek (1977) process, which suggests the extraction of energy and angular momentum from the spinning black hole.   In this paper, we assume a stationary and axisymmetric magnetosphere, and consider ideal MHD flows along a magnetic field line. The initial velocity can be at most less than the slow magnetosonic wave speed. To accrete onto the black hole, the ejected inflows from the the plasma source must pass through the slow magnetosonic point ({\\sf S}), the Alfv\\'en point ({\\sf A}) and the fast magnetosonic point ({\\sf F}) in this order, as it is well known. At these points, {\\sf A}, {\\sf F} and {\\sf S}, the poloidal velocity equals one of the Alfv\\'{e}n wave and fast and slow magnetosonic wave speeds, respectively. In the case of accretion onto a star, because the accreting plasma is stopped at the stellar surface, a shock front would be formed somewhere on the way to the stellar surface and the accretion becomes sub-fast magnetosonic. However, for accretion onto a black hole, the flow must be super-fast magnetosonic at the event horizon ({\\sf H}). If not so, the fast magnetosonic wave can extract information from the interior of the black hole to the exterior; this fact obviously contradicts with the definition of the event horizon. In fact, an ideal MHD accretion solution which keeps sub-fast magnetosonic has zero poloidal velocity at the event horizon and the density of the plasma diverges; the solution is unphysical.   Because the magnetic field lines would rigidly rotate under the ideal MHD assumption, there are two light surfaces ({\\sf L}) in the black hole magnetosphere \\citep[hereafter Paper~I]{Znajek77,Takahashi-NTT90}. The plasma source must be located between these two surfaces. Further, one or two Alfv\\'en surfaces lie between the two light surfaces, and for accretion there must be a fast-magnetosonic surface between the Alfv\\'en surface and the event horizon (see Paper~I). Here, we should note that the physical mechanism to determine the angular velocity of the field lines is controversial. A time-dependent determination of it has been discussed by \\citet{Punsly01}; the torsional Alfv\\'en wave originated from the plasma source and propagated up and down the magnetic flux tube forces to minimize the magnetic stresses in the system.   The conditions on the flows at the magnetosonic points and the Alfv\\'en point restrict the five physical parameters which specify the flow (see the following section) if one is to adhere to the ideal MHD assumption globally. The fast and slow critical points have X-type (physical) or O-type (unphysical) topology for the solution, while the Alfv\\'{e}n point does not specify its topological feature; hereafter, we will call the flow passing through both X-type fast and X-type slow magnetosonic points as the ``{\\sf SAF}-solution''. When we discuss the global features of a solution, it is very important to know the numbers of these critical points and the Alfv\\'en points. In a zero-temperature limit (cold limit), the regularity condition for the trans-fast MHD flow was discussed by \\citet{Takahashi94}. In this case, the Alfv\\'en points and the fast magnetosonic points only appear in wind and accretion solutions without the slow magnetosonic point, because the velocity of the slow magnetosonic wave speed is zero. The relativistic hot MHD flow equation has been formulated by \\citet{Camenzind86a,Camenzind86b,Camenzind87,Camenzind89}; see also, Paper~I. In \\S~2, we summarize the basic equations for the MHD flows and the condition at the Alfv\\'en point discussed in Paper~I.   Though we consider the general relativistic plasma flow, the significance of the Alfv\\'en point and the fast and slow magnetosonic point conditions is similar to that of a Newtonian wind model by \\citet{Weber-Davis67} and a special relativistic wind model by \\citet{Kennel83}. \\citet{Kennel83} classified the outgoing trans-Alfv\\'enic MHD wind solutions into a ``critical'' ($E=E_{\\rm F}$) solution, ``sub-critical'' ($E<E_{\\rm F}$) solutions and ``super-critical'' ($E>E_{\\rm F}$) solutions, where $E$ is the conserved energy of the wind and $E_{\\rm F}$ is the energy for the trans-fast MHD wind. To reach distant regions, the critical (trans-fast MHD wind) solution and the super-critical (sub-fast MHD winds) solutions are physical, and the sub-critical solutions are unphysical beyond the turnaround point. Turning now to the ingoing plasma flow, the topology of the black hole accretion solution space also has a similar structure; that is, (i) a trans-Alfv\\'en MHD ingoing flow with $E = E_{\\rm F}$, which means a trans-fast MHD ingoing flow discussed in this paper (critical), (ii) trans-Alfv\\'en MHD ingoing flows with $E < E_{\\rm F}$ (sub-critical) and (iii) trans-Alfv\\'en MHD ingoing flows with $E > E_{\\rm F}$ (super-critical). Under the ideal MHD approximation, the physical solution is only the critical solution (i); the super-critical solutions (iii) are unphysical for the reason mentioned above. In addition to these, for accretion onto a black hole, we must consider (iv) sub-Alfv\\'en (or sub-slow MHD) ingoing flows, although they do not pass through the Alfv\\'en point {\\sf A}. The breakdown of ideal MHD approximation between the horizon and the inner light surface is indicated by \\citet{Punsly01}, and then non-ideal MHD solutions classified into (ii), (iii) and (iv) would be realized as accretion solutions onto a black hole (see \\S~5).   The main purpose of this paper is to examine the thermal effects on an ideal MHD plasma streaming in a black hole magnetosphere (see Fig.~\\ref{fig:acc}) by studying the critical conditions at those magnetosonic points. Now, the slow magnetosonic point appears on the MHD flow solutions. The details of critical conditions at the fast and slow magnetosonic points are discussed in \\S~3. We derive the critical conditions at the fast and slow magnetosonic points, which are denoted in terms of the location of the fast and slow magnetosonic points, the sound velocity at the fast and slow magnetosonic points and the locations of the Alfv\\'{e}n and light surfaces. In \\S~4, we clarify the thermal effects on the MHD flows, and discuss its dependence on the rotation of the black hole magnetosphere and the divergence of the cross-section of a magnetic flux-tube along the field line. Then, we can find two kinds of trans-fast MHD flow solutions for both inflows and outflows: ``hydro-like'' MHD flow and ``magneto-like'' MHD flow. The main difference between the two solutions is the behavior of the magnetization parameter, which is the ratio of the fluid and electromagnetic parts of the total energy of the flow. The hydro-like MHD flow solution is a somewhat hydrodynamical solution and, in the weak magnetic field limit, this trans-fast magnetosonic flow solution becomes a trans-sonic flow solution discussed by \\citet{Abramowicz81} and \\citet{Lu86} in the hydrodynamical case. We also unify hydrodynamic flows with hot MHD flows in a common formalism. The hydro-like MHD flow solution disappears for a magnetically-dominated magnetosphere. On the contrary, the magneto-like MHD flow solution results in the magnetically-dominated flow, although that disappears for hotter plasma cases. In \\S~5, we summarize our results.   \\placefigure{fig:acc}% ",
        "conclusions": "%  We have considered stationary and axisymmetric hot ideal MHD accretion along a flux-tube connected from a plasma source to the event horizon. To argue the details of the boundary conditions at the plasma source would take us beyond the scope of this paper. Therefore, we have surveyed the dependence of the trans-fast MHD flows on a wide range of source parameters. We have shown that, when the forbidden region is type IA or~IIIA, there exist two physically different accretion regimes: (i) {\\it magneto-like}\\, MHD accretion and (ii) {\\it hydro-like}\\, MHD accretion. The magneto-like MHD accretion would be injected from near the separation point and passes through the inner Alfv\\'en point with a smaller $\\hat\\eta$. On the other hand, the hydro-like MHD accretion with a sufficiently large $\\hat\\eta$ would be injected from between the outer Alfv\\'en point and the outer light surface and passes through the outer Alfv\\'en point. Hydro-like accretion may also be initially super-slow magnetosonic or super-Alfv\\'enic. A hot ideal MHD plasma with larger $\\hat\\eta$ cannot accrete stationary onto the black hole after passing through the inner Alfv\\'en point. Then, if the value of $\\hat\\eta$ increases with a secular timescale, the magneto-like MHD accretion solution should transit to the hydro-like MHD accretion solution; the inverse process would be also possible. The criterion for distinguishing between the two regimes is based on the locations of both the Alfv\\'en point and the fast magnetosonic point, which change discontinuously during the transition.   For the magneto-like MHD accretion, we have found that the location of the X-type fast magnetosonic point is not unique for fixed intrinsic parameters of the accreting plasma. For example, in Figures~\\ref{fig:eta-OM}a,~\\ref{fig:eta-spin}b and~\\ref{fig:eta-alf}, in the range of $r_{\\rm H}<r_{\\rm F}<r_{\\rm A}^{in}$, a $\\hat\\eta=$ constant line crosses a solid curve of $\\zeta_{\\rm F}\\leq 0.1$ at three points; the first and third fast magnetosonic points are X-type critical points, while the second one is O-type. So, we can expect two modes of magneto-like MHD accretion solutions; that is, one passes through the ``inner'' inner-fast magnetosonic point and the other passes through the ``outer'' inner-fast magnetosonic point. The number of these inner-fast magnetosonic points depends effectively on $a$, $\\delta$, $\\Omega_F$, $\\tilde L$, $\\zeta_{\\rm F}$ and $\\hat\\eta$. There is a tendency for larger $\\Omega_F-\\omega_{\\rm H}$ to generate a multiple inner-fast magnetosonic solution. Compared with the radial field geometry, converging field geometries for accreting flows ($\\delta>0$) also generate such multiple inner-fast magnetosonic solutions. Transition between these two modes is also discontinuous. The $\\delta$-dependence of {\\sf SAF}-accretion solutions on the flow velocity $u^r(r)$ and the electromagnetic energy $X_{\\rm em}(r)$ is very weak, as long as the location of the fast magnetosonic point does not jump, by changing the $\\delta$ value, to another branch of the multi-inner fast magnetosonic points. The radial terminal velocity at the event horizon $u^r_{\\rm H}$, rather than the radial one, slightly decreases (increases) for $\\delta>0$ ($\\delta<0$). For example, we have checked this for $\\delta=0.4, 0.0$ and $-0.4$ cases. For such accretion solutions, we also find that for $\\delta>0$ ($\\delta<0$) the total energy flux per magnetic tube $\\hat\\eta E$ and $\\hat\\eta$ increase (decrease), while the total energy $E$ decreases (increases).   Throughout this paper, we have only discussed the ideal MHD flow cases. However, non-ideal MHD flow solutions near the event horizon are also presented (\\cite{Punsly90,Punsly01}). \\cite{Punsly90} discussed an ingoing magnetic flow solution along magnetic field lines that thread the ergosphere and the equatorial plane (and therefore not the event horizon). This solution corresponds to our sub-Alfv\\'enic ingoing solution approaching the inner light surface with zero poloidal velocity (see, e.g., Fig.~\\ref{fig:MHDaccA}a); along this solution $B^\\phi=0$ at $r=r_{\\rm A}^{in}$ (see also, Fig.~\\ref{fig:MHDaccA}b), where $u_p^2\\neq u_{\\rm AW}^2$ (not the Alfv\\'en point {\\sf A}), while for the {\\sf SAF}-solution $B^\\phi\\neq 0$ at the Alfv\\'en point {\\sf A}. When we consider an accretion solution onto a black hole, it seems that there is no sub-Alfv\\'en accretion solution consistent with the ideal MHD approximations due to the existence of the forbidden region. However, of course, for such a set of ingoing flow parameters the ideal MHD approximation must be rejected, and then the non-ideal MHD ingoing flow should exist. This is because near the light surface due to the plasma inertia effects large radiation losses are expected, and large radiation losses equate to a dissipative plasma and a breakdown of the ideal MHD approximation. Then, a non-ideal MHD accretion flow solution, which does not pass through the fast-magnetosonic critical point {\\sf F} discussed in the previous section, also exists in the region downstream of the light surface because of the inward attraction of black hole gravity. This requires that dissipative effects be incorporated into the physical description of the inward extension of the ideal MHD wind inside the light surface. Such a plasma propagates inside the inner light surface and enters the forbidden region of $r\\leq r_{\\rm L}^{in}$ with relatively slow velocity (as compared with the Alfv\\'en wave speed except for the area close to the horizon) by crossing or reconnecting the magnetic field lines, where the physical meaning of the concept of the forbidden region for the ideal MHD flows would be lost. Note that the non-ideal MHD accretion flow must also become super-Alfv\\'en and super-fast magnetosonic just inside the inner light surface (see Chapter~9 of Punsly 2001).   For a ``super-critical'' accretion flow of $E>E_{\\rm F}$ and $u_{\\rm AW}^2 < u_p^2 < u_{\\rm FM}^2$, which reaches the event horizon without passing through the fast magnetosonic point {\\sf F} and is unphysical under the ideal MHD approximation, some kinds of dissipative effects near the event horizon would also make the super-critical accretion onto the black hole possible. One may also expect a ``sub-critical'' accretion flow of $E<E_{\\rm F}$, which is also not a solution for the inflow into the black hole under the ideal MHD approximation. However, the breakdown of the ideal MHD would change the nature of the critical points (see, e.g., \\cite{Chakrabarti90} for the hydrodynamical flow case); that is, instead of the X-type and O-type critical points, the ``nodal-type'' and ``spiral-type'' critical points appear on the flow solutions in the ($r, u^r$)-plane. Thus, when the initial condition at the plasma source does not match the critical condition at the fast magnetosonic point, the non-ideal MHD inflow into the black hole must be realized. The detailed structure of the non-ideal MHD flows around the critical point is a very important topic for black hole accretion, but further discussion is out of the scope of this paper.   When the angular velocity of the magnetosphere is in the range of $0<\\Omega_F<\\omega_{\\rm H}$, only one Alfv\\'en point appears on the $r$-$u^r$ plane. To accrete onto the black hole, an ideal MHD flow must pass through the Alfv\\'en point classified by type IIA$^{in}$, IIA$^{out}$, IIB$^{in}$ or IIB$^{in}_\\ast$ in Paper~I. Determining whether the accretion is magneto-like or hydro-like may not be clear in type IIA case. However, we can expect that the Alfv\\'en point {\\sf A}$^{in}$ is responsible for the magneto-like accretion solution and the {\\sf A}$^{out}$ point is responsible for the hydro-like accretion solution. It seems that for the Alfv\\'en point of type IIB$^{out}$ there are no accretion solutions consistent with the ideal MHD approximation. Of course, for such a set of ingoing flow parameters the ideal MHD approximation must be abandoned. When the {\\sf SAF}-accretion solution is invalid (e.g., the type~IIB$^{out}$ case), the non-ideal MHD accretion becomes important and would dynamically effect the magnetic field structure.   When a region of $k(r) > 0$ exists along a magnetic field line for larger $\\tilde L\\Omega_F \\, (\\leq 1)$, the forbidden region is type~B and the ideal MHD accretion is only allowed after passing the {\\it inner}\\/ Alfv\\'en point. The hydro-like MHD accretion does not arise because of the sufficiently strong centrifugal barrier, while the magneto-like MHD accretion is available because of the effective angular momentum transport from the fluid-part of total angular momentum to the magnetic-part (Hirotani et~al.~1992). When the hot effects dominate in the plasma, this magneto-like MHD accretion may be also forbidden due to the disappearance of the inner-fast magnetosonic point. In this case, however, we can also expect the non-ideal MHD accretion to fall into the black hole.   If a shock front is generated after passing the fast magnetosonic point, the post-shock flow with increased entropy must pass another fast magnetosonic point again on the way to the event horizon. To construct such a shock formation model, the existence of multiple fast magnetosonic points is required in the accretion solution. We can expect two types of discontinuous transitions. One is the transition from the hydro-like solution to the magneto-like solution at somewhere between the middle-fast and inner-fast magnetosonic points. The other is the transition in the hydro-like or magneto-like solution. For the magneto-like solution, we can see the possibility of transition from the $r_{\\rm cr}$-$\\hat\\eta$ diagrams. That is, the accreting matter passes through the first inner-fast magnetosonic point located just inside the inner-Alfv\\'en point, and then after the shock it passes through the third (X-type) inner-fast magnetosonic point. Here, the first and third inner-fast magnetosonic points are located on different $\\zeta_{\\rm F}$-curves in the $r_{\\rm cr}$-$\\hat\\eta$ diagram because of the entropy generation. The insights gained in the course of our analysis should be of use in further investigations of shocked accretion solutions.   Although we have discussed accretion flows onto a black hole, our treatment can be applied to outgoing flows (i.e., winds and jets). To do so, we can plot $r_{\\rm cr}$-$\\hat\\eta$ diagrams in the range of $r_{\\rm L}^{in} < r_{\\rm cr} < \\infty$ (for the fast magnetosonic point, $r_{\\rm A}^{in} < r_{\\rm F} < \\infty$; for the slow magnetosonic point, $r_{\\rm L}^{in} < r_{\\rm S} < r_{\\rm A}^{out}$). So, we will find the possible locations of the magnetosonic point for outflows. When the forbidden region is type IA or IIIA, we find two regimes of {\\sf SAF}-solutions for outflows; that is, ${\\sf inj}$ $\\to$ ${\\sf S}^{\\rm mid}$ $\\to$ ${\\sf A}^{out}$ $\\to$ ${\\sf F}^{\\rm out}$ $\\to$ $\\infty$ and ${\\sf inj}$ $\\to$ ${\\sf S}^{\\rm in}$ $\\to $ ${\\sf A}^{in}$ $\\to$ ${\\sf F}^{\\rm mid}$ $\\to$ $\\infty$, where the slow magnetosonic point ${\\sf S}^{\\rm in}$ is located between the inner light surface and the inner Alfv\\'en point, and the fast magnetosonic point ${\\sf F}^{\\rm out}$ is located outside the outer Alfv\\'en point. The former {\\sf SAF}-solution is also available for the type B forbidden region, and the latter {\\sf SAF}-solution remains in the hydrodynamical limit. From the $r_{\\rm cr}$-$\\hat\\eta$ diagrams, for the outflow, the location of the fast magnetosonic points with the same $\\hat\\eta$ value strongly depends on $\\delta$. A similar result has been discussed in \\citet{Takahashi-Shibata98} as a pulsar wind model without the gravitational effects. Here, we should note that to blow away to infinity, the outgoing flows also must pass through the slow magnetosonic point and the Alfv\\'en point, but at the asymptotic region both the super-fast and sub-fast magnetosonic outflows are available. Whether or not the outflow passes through the fast magnetosonic point depends on the field aligned parameters of flows and the geometry of field lines. So, for the outflow, the fast-magnetosonic surface does not need to cover all solid angles. We can expect that the trans-fast MHD outflow is realized at least in some part of the magnetic field lines, and the distribution of the fast-magnetosonic surface would play a very important part in explaining the generation and collimation of a highly accelerated jet or wind."
    },
    "0201/astro-ph0201057_arXiv.txt": {
        "abstract": "{The gravitational potential responsible for the lensing effect in \\SBS\\, is studied over length scales from a few arc-seconds to a few arc-minutes. For this purpose, we use sharply deconvolved Hubble Space Telescope images in the optical and near-IR, in combination with ground based optical data obtained over a wider field-of-view. In particular, we have carried out a multi-color analysis in order to identify groups or clusters of galaxies along the line of sight. Photometric redshifts are measured for 139 galaxies unveiling significant excesses of galaxies  1.0\\arcmin\\, NW and 1.7\\arcmin\\, SW of the main lensing galaxy. The photometric redshift inferred both for the main lensing galaxy and for the galaxy concentrations is z=$0.9^{+0.10}_{-0.25}$. This is in rough agreement with the measured spectroscopic redshift of the main lensing galaxy, z=$0.71$ (Burud et al. 2002), suggesting that it is part of a larger group or cluster. We investigate the impact of including the galaxy cluster, first on the modelling of the lensing system, and second on the expected time--delay between the two quasar images. ",
        "introduction": "The study of multiply imaged quasars is one of the most promising way to measure cosmological parameters such as the Hubble parameter \\ho\\, (Refsdal 1964, Blandford \\& Narayan 1992). Indeed, \\ho\\, is related to two observables: the time--delay between the light curves of the quasar images and the mass distribution in the lens.  Until recently, simple mass distributions were used to model multiply imaged quasars. However, the improvement of observing techniques has led to the discovery that most image configurations require more complex models, involving, for example, an external shear (Keeton, Kochanek \\& Seljak 1997).  This small but significant external perturbation to the main lensing potential often corresponds to the presence of groups or even clusters of galaxies along the line of sight (see for example Keeton \\& Kochanek 1997, Burud et al. 1998, Morgan et al. 2001, Fassnacht \\& Lubin 2001). These additional structures must be taken into account in order to properly model the mass distribution and accurately convert the observed time--delay into \\ho.\\\\ \\SBS\\, is a doubly imaged quasar at z=$1.855$, with an angular separation of 1.568\\arcsec\\, ($\\alpha_{2000}$=15$h$\\,21$m$\\,44.83$s$\\,, $\\delta_{2000}$=+52$\\degr$\\,54\\arcmin\\,48.6\\arcsec). It was discovered by Chavushyan et al. (1997) in the course of the Second Byurakan Survey (Markarian \\& Stepanian, 1983). Crampton et al.  (1998) identified the lensing galaxy from near-IR adaptive optics images and proposed a simple lens model including exclusively the main lensing galaxy.  They concluded however that an additional external shear was mandatory to model properly the system.\\\\ Deep ground--based and HST data are used in the present paper to map the mass distribution along the line of sight to SBS 1520+530. We then investigate the source of external shear and propose a multi--components model for the total lensing potential, where the main lensing galaxy is a member of a larger group or cluster of galaxies.   \\begin{figure*}[hbtp] \\begin{center} \\includegraphics[width=18cm]{faure1.ps} \\caption{\\label{fig1} Simultaneous deconvolution of the four HST/NICMOS2 frames obtained of \\SBS. The left panel shows a standard combination of the data. The right panel is the result of the deconvolution, with a final resolution of 0.075\\arcsec. The pixel size on this 19\\arcsec wide image is 0.0375\\arcsec. North is to the top, East is to the left. The labels follow the description in section \\ref{result}.} \\end{center} \\end{figure*} ",
        "conclusions": "\\label{sec4}  We have studied the doubly imaged quasar \\SBS\\, using both HST and ground based data, in order to map the projected mass distribution of the lensing system at small and large scales.  First, by deconvolving the HST/NICMOS data, we have accurately measured the shape and brightness of the main lensing galaxy. We detect a weak extension to the NW of the lensing galaxy (Fig.~1).  This feature does not seem to be a residual of the deconvolution process. If it is a genuine feature, it may have some influence on the lens model, considering its proximity to image A of the quasar. However until it has been confirmed, we ignore it for the modelling.   Second, we have mapped the relative surface density of faint galaxies around the quasar images and identified an overdensity of objects to the West. Photometric redshifts were estimated for the main lensing galaxy L as well as for all the galaxies in the 7.5~(\\arcmin)$^2$ surrounding the lensed quasar. This analysis reveals a concentration in redshift space around a mean value z=0.9$^{+0.10}_{-0.25}$ and confirms the 3D--reality of the observed projected overdensity of objects: we consider it to be a galaxy cluster. As a consequence, it is probable that the main lensing galaxy L (measured spectroscopic redshift z=0.71) is a member of this cluster (measured photometric redshift z=0.9$^{+0.10}_{-0.25}$). Following this argument and considering that the spectroscopic redshift measurement is more accurate, we have built our lens model with a unique redshift value,  z=0.71, for the main lensing galaxy, for galaxy M and for the cluster.    Finally, the presence of a galaxy cluster to the W of \\SBS\\, is also suggested by the lensing model analysis which does require such a contribution. If not taken into account, it is impossible to reproduce the observed configuration and flux ratio of the double quasar images. This argument is valid as long as light traces mass, which seems to be the case for most lensing galaxies (Kochanek 2001).   An independent and complementary way of probing the line of sight mass distribution is to measure the time--delay induced by the lens system (also necessary to constrain \\ho). Therefore, it is interesting to predict the time-delay corresponding to the lens model derived from the lensed quasar image configuration. We have made time-delay predictions for the 3 different lens models analysed in this paper, to be compared with the observed time-delay, when it will be available.   Yet, further detailed observations are needed to improve the understanding of the lensing system and enable to use it as a cosmological probe.   \\SBS\\, will be observed by the X--ray satellite {\\it Chandra} (PI: G.Garmire): this observation should confirm the existence of the cluster in a way similar to the quadruply lensed quasar RX~J0911.4+0551 (Morgan et al. 2001). Deep Keck/Gemini multi-object spectroscopy would be invaluable to settle the redshift of the detected cluster/group and provide an estimate of its mass."
    },
    "0201/astro-ph0201261_arXiv.txt": {
        "abstract": "We present the results of fully 3-D hydrodynamic simulations of the gravitational collapse of isolated, turbulent molecular cloud cores. Starting from initial states of hydrostatic equilibrium, we follow the collapse of both singular and nonsingular logatropic cores until the central protostar has accreted $> 90\\%$ of the total available mass.  We find that, in the collapse of a singular core with access to a finite mass reservoir, the mass of the central protostar increases as $M_{\\rm acc} \\propto t^{4}$ until it has accreted $\\sim 35\\%$ of the total available mass.  For nonsingular cores of fiducial masses 1, 2.5, and 5 M$_{\\odot}$, we find that protostellar accretion proceeds slowly prior to the formation of a singular density profile.  Immediately thereafter, the accretion rate in each case increases to $\\sim 10^{-6}$ M$_{\\odot}$ yr$^{-1}$, for cores with central temperature $T_{c}= 10$ K and truncation pressure $P_{s} = 1.3 \\times 10^{5} k_{B}$ cm$^{-3}$ K.  It remains at that level until half the available mass has been accreted.  After this point, the accretion rate falls steadily as the remaining material is accreted onto the growing protostellar core.  We suggest that this general behaviour of the protostellar accretion rate may be indicative of evolution from the Class~0 to the Class~I protostellar phase. ",
        "introduction": "A crucial prerequisite for any comprehensive theory of isolated star formation is an understanding of the gravitational collapse of molecular cloud cores\\footnote{To minimize ambiguity, we will use the terms ``cloud'', ``clump'', and ``core'' in the following hierarchical manner:  molecular \\emph{clouds} contain dense condensations called \\emph{clumps}, in which entire star clusters may form; clumps, in turn, contain \\emph{cores}, from which single or binary stars may form.  The centre of a core shall be denoted herein with the terms ``central object'' or ``central protostar''.}.  \\citet{Lar69} and \\citet{Pen69} pioneered the study of the collapse of the isothermal sphere from both analytical and numerical perspectives.  Subsequently, \\citet{Shu77} defined the long-standing paradigm for isolated low mass star formation---the ``inside-out'' or ``expansion wave'' collapse---with his elegant self-similar solution for the collapse of a singular isothermal sphere (SIS).  A common feature of all three of the aforementioned studies, and most star-formation studies which have followed, is their adoption of the isothermal equation of state (EOS).  Recent evidence supports the assertion that some observed molecular cores are well-fit by isothermal models.  For example, \\citet{Alves} demonstrated convincingly that the density distribution of the dark cloud Barnard 68 is very well fit by a critically stable Bonnor-Ebert sphere \\citep{B56,E55}.  There is, however, mounting evidence which indicates that more massive molecular cores possess properties which cannot be explained by the isothermal EOS.  Surveys of prestellar cores in a variety of star-forming complexes show that massive cores have strong nonthermal linewidths which are greater than their thermal linewidths \\citep{CM95}. Furthermore, it has been shown by \\citet{FM92} that even low mass cores have significant nonthermal linewidths.  The SIS model for molecular cloud core collapse predicts a time-invariant rate of accretion onto the central protostar \\citep{Shu77}, which gives rise to a large associated spread between the formation times of high and low mass stars.  In the high-mass star formation regime, this is problematic: the constant accretion rate predicted by SIS models is too low to form massive stars on the timescales suggested by observations (\\citeauthor{CM95} \\citeyear{CM95}; \\citeauthor{MP97} \\citeyear{MP97}, MP97 hereafter).  \\citet{FC93} and \\citet{BB} have shown that it is possible to obtain variable accretion rates from isothermal collapse models if, for example, the mass reservoir is finite or there are initial deviations from the SIS density profile.  It is not clear, however, that the mass accretion rates generated by these models vary in a way which is consistent with observations.  The likely conclusion is that isothermal models only apply to the formation of low mass stars.  The SIS collapse model assumes hydrostatic initial conditions; it has been suggested that such a singular configuration could arise through the subsonic collapse of a Bonnor-Ebert sphere \\citep{SAL87}.  Simulations of the collapse of critical Bonnor-Ebert spheres, by contrast, indicate that, as the density achieves a singular form (i.e. when $\\rho(r)$ takes on an $r^{-2}$ profile), 44\\% of the mass is infalling at a few times the sound speed (\\citeauthor{FC93} \\citeyear{FC93}; FC93 hereafter).  Such high infall velocities have not yet been detected in collapsing cores, but whether this is because they do not exist, or because their detection is beyond the limits of current observational techniques is not yet known (see FC93 for discussion).  One proposed alternative to the isothermal EOS is the \\emph{logatropic} equation of state, so named for its logarithmic dependence of pressure on density.  First proposed by \\citet{LS89}, the logatrope is an empirically-motivated EOS which attempts to account for the presence of nonthermal or turbulent pressure within molecular cores. In the so-called ``mixed'' form of \\citet{LS89}, the logatrope was essentially just the isothermal equation of state with a correction added to account for turbulent support.  \\citet{MP96} (MP96 hereafter) proposed the ``pure'' form of the logatrope, which eliminated the linear (isothermal) dependence of pressure on density, leaving only the logarithmic dependence.  In this form, the logatrope can describe the properties of both low and high mass cores (we refer here primarily to the properties of molecular cores as determined by spatially unresolved measurements, such as single measurements of the thermal and nonthermal components of a core's velocity dispersion).  While nonthermal support is certainly important in high mass cores, low mass cores are also characterized by some degree of nonthermal support.  The logatrope is capable of addressing some of the difficulties with the isothermal models.  For example, it is specifically formulated to account for the nonthermal linewidths of observed cores.  Also, because they predict a variable protostellar mass accretion rate ($\\dot{M}_{acc} \\propto t^{3}$), self-similar models for the collapse of a singular logatropic sphere (SLS) give rise to a relatively small spread between the formation times of high and low mass stars.  The general consensus at this time favours models which produce subsonic molecular cloud core collapse.  Because the logatrope is a softer EOS than the isothermal EOS (i.e. $P$ depends less than linearly on $\\rho$), it has been suggested that it might allow for a gentler, subsonic approach to the singular density profile (MP96).  This contrasts with the strongly supersonic flow found in simulations of the collapse of Bonnor-Ebert spheres (FC93).  The logatrope has the potential to produce gentler collapses because it accounts for some measure of turbulent and magnetic support within a molecular cloud core.  There is direct observational evidence that the internal (i.e. spatially resolved) structure of individual turbulent cores matches the predictions of logatropic models.  For example, \\citet{vdT} examined the structure of the envelopes around 14 massive young stars and found that the density profiles of these envelopes have power law distributions, $n \\propto r^{\\alpha}$, with indices $\\alpha = 1.0-1.5$.  This is consistent with the logatropic model, which predicts $\\alpha = 1.0$ for cores in equilibrium and $\\alpha =1.5$ for collapsing cores.  Further, these values of $\\alpha$ indicate density profiles which are significantly flatter than the $\\alpha = 2.0 \\pm 0.3$ commonly found for low mass cores, and are inconsistent with the value of $\\alpha = 2.0$ predicted by the SIS model for an equilibrium isothermal core. These results indicate that nonthermal support mechanisms dominate in massive young stellar objects, while thermal pressure dominates in objects of lower mass.  Similarly, \\citet{Henning} and \\citet{Colome} find $\\alpha = 0.75-1.5$ for the envelopes around massive Herbig Ae/Be stars.  Other evidence is provided by \\citet{OLA}, who examined the dust thermal spectra of several hot molecular cores, which are thought to be the sites of massive star formation.  They showed that the thermal spectra of these cores were best fit by models whose envelopes had the density profile of collapsing logatropic spheres.  The equilibria, stability criteria, and self-similar collapse solutions for pure logatropic spheres have been derived by MP96 and MP97. They showed that logatropic spheres have two equilibria: one is a stable equilibrium with finite central density and a $\\rho \\propto r^{-1}$ envelope, and the other is an unstable equilibrium with a singular ($\\rho \\propto r^{-1}$) density profile throughout.  These equilibria are roughly analogous to those of the isothermal sphere, namely the hydrostatic Bonnor-Ebert sphere and the unstable SIS.  MP97 also derived self-similar collapse solutions for the SLS, similar to those developed by \\citet{Shu77} for the SIS.  In this paper, we extend the work of MP97 by performing fully three-dimensional numerical hydrodynamic simulations of the gravitational collapse of logatropic spheres, both singular and nonsingular.  Our goal is to develop a set of numerical collapse solutions for the logatrope which is analogous to those developed for the Bonnor-Ebert sphere by FC93. We present results of simulations which describe the approach to a singular density profile in the collapse of an initially hydrostatic, nonsingular logatropic sphere.  Because the collapse of a nonsingular logatrope is not suited to exploration by the usual self-similar analytical techniques, we are motivated to use a more flexible numerical approach.  In \\S \\ref{sec:lc}, we review the analytic theory of logatropic equilibria and collapses.  Section \\ref{sec:nm} describes our numerical method and discusses our testing procedures. Sections \\ref{sec:singcol} and \\ref{sec:nscol} present the results of our simulations of the singular and nonsingular logatropic collapses (scaled such that the total core mass is 1 M$_{\\odot}$), respectively.  Section \\ref{sec:obsmatch} discusses the results of simulations of the collapse of higher-mass nonsingular cores and relates them to the observed properties of protostellar cores.  We conclude with discussion and a summary in sections \\ref{sec:disc} and \\ref{sec:summ}. ",
        "conclusions": "\\label{sec:disc}  As star-forming cores evolve from Class 0 to Class IV, their protostellar mass accretion rate may vary by $\\sim 2$ orders of magnitude (see \\citeauthor{AndrePPIV} \\citeyear{AndrePPIV}, and references therein). It should be noted that estimates of the accretion rate in young stellar objects are not made by direct measure, but rather by inference from a combination of models and other data, as will be discussed shortly.  In contrast to the observations, the standard SIS model predicts a time-invariant protostellar mass accretion rate.  \\citet{FC93} and \\citet{BB} have obtained variable mass accretion rates for isothermal collapse models which draw mass from finite reservoirs, but these models show a tendency to evolve toward the constant mass accretion rate predicted by the SIS model.  \\citet{Henriksen} and \\citet{CM95} have proposed alternative models for protostellar collapse which predict variable accretion rates.  Like the logatrope, these models are phenomenologically based, designed to match the observed properties of star-forming cores (density profiles, nonthermal linewidths, etc.) These models, however, have many more free parameters than the logatrope, which has only one: the constant $A$ in equation (\\ref{eq:eqos1}).  As shown in the previous section, the logatrope also predicts a highly variable accretion rate.  We must ask, then, how closely the logatropic accretion profile corresponds to available observational data.  \\citet{Bontemps} surveyed 36 Class~I and 9 Class~0 protostars in the $\\rho$~Oph, Taurus-Auriga, and Perseus star-forming regions and found that the outflow momentum flux, $F_{\\rm CO}$ declines from $\\sim 10^{-4}$ M$_{\\odot}$ km s$^{-1}$yr$^{-1}$ in the Class~0 sources to $\\sim 2 \\times 10^{-6}$ M$_{\\odot}$ km s$^{-1}$yr$^{-1}$ in the Class I sources.  Their suggestion was that this decrease in $F_{\\rm CO}$ and the accompanying decrease in the mass ejection rate of the protostellar wind, $\\dot{M}_{w}$, were attributable to a decrease in the protostellar mass accretion rate, $\\dot{M}_{\\rm acc}$, from $\\sim 10^{-5}$ M$_{\\odot}$ yr$^{-1}$ to $\\sim 2 \\times 10^{-7}$ M$_{\\odot}$ yr$^{-1}$. Clearly the largest of these accretion rates is significantly higher than those shown in Figure \\ref{fig:mdot}, but we note that there is significant room for movement in both the observations and the scaling of our simulations.  We have scaled our accretion profiles to central temperatures and truncation pressures which are characteristic of regions of isolated star formation.  If we instead choose $T_{c}$ and $P_{s}$ in keeping with observations of regions such as $\\rho$ Oph, we can make up the difference between our $\\dot{M}_{\\rm acc}$ values and those suggested by \\citet{AndrePPIV}.  \\citet{Johnstone} fitted 55 molecular cloud cores in $\\rho$ Oph to Bonnor-Ebert spheres and thereby made estimates of the surface pressure on each core.  The typical truncation pressures found in that study lay in the range $P_{s} = 10^{6}-10^{7} k_{B}$ cm$^{-3}$ K. These are between 1 and 2 orders of magnitude greater than the fiducial truncation pressure used herein ($ P_{s} = 1.3 \\times 10^{5} k_{B}$ cm$^{3}$ K).  The relationship between the total mass of a nonsingular logatropic core and its truncation pressure takes the form $M_{\\rm tot} \\propto P_{s}^{-1/2}$.  Hence, increasing the truncation pressure by a factor of 100 would reduce the mass of a given core by a factor of 10. This would also bring the mass of the critical $A=0.2$ logatrope down from 92 M$_{\\odot}$---typical of an entire molecular cloud \\emph{clump}---to 9.2 M$_{\\odot}$, which is typical of a molecular cloud \\emph{core} situated within a clump.  Hence, our models would then represent cores of comparatively low mass.  We would then have to look at logatropic spheres closer to the critical radius in order to find cores of mass $\\sim 1$ M$_{\\odot}$, in which case the accretion rates ought to be considerably higher than those listed in Table \\ref{tab:nstable}.  The method used by \\citet{Johnstone} to extract truncation pressures was, however, dependent on the assumption that all of the cores could be well-fit by Bonnor-Ebert spheres, and hence the truncation pressures extracted cannot necessarily be considered reflective of those that would be required to truncate logatropic spheres.  The dimensionless results of Figures \\ref{fig:mdot}a and \\ref{fig:mdot}b are meant to be rescalable, should future EOS-independent measures of $P_{s}$ become available.  Disregarding the absolute numbers for a moment (which are, as we have emphasized, subject to considerable uncertainty), we note that the \\emph{trend} in our logatropic $\\dot{M}_{\\rm acc}$ bears significant resemblances to the data.  Estimates place the transition from the Class~0 to the Class~I protostellar phase at the point when the accreted mass is approximately equal to that remaining in the protostellar envelope, $M_{\\rm acc} \\simeq M_{\\rm env}$ \\citep{AndrePPIV}.  As shown in Table \\ref{tab:mdotab}, the nonsingular logatropic collapse behaves in just this way---once about half of the mass of the core has been accreted, the accretion rate, $\\dot{M}_{acc}$, begins to fall off.  This behaviour arises because the accreting central object has only a finite mass reservoir from which to draw (see discussion in \\S~\\ref{sec:singcol}).  We suggest that one interpretation of the mass accretion profiles in Figure \\ref{fig:mdot} is that they describe the transition from Class~0 to Class~I objects.  If we take the 1 $M_{\\odot}$ case as an example, we see that, at times $t < 0$, there is little accretion activity, so this may correspond to a very early Class~0 or even a preprotostellar stage.  From $t \\sim 0 $ to $\\sim 0.7 \\times 10^{6}$ yr, at which time the mass of the protostar and the remaining envelope are equal, the object undergoes a period of vigorous accretion in an SLS-like manner.  During this interval, the accretion rate, while variable, stays relatively constant, never deviating by more than $6\\%$ from the average of $7.4 \\times 10^{-7}$ M$_{\\odot}$ yr$^{-1}$.  If the relationship between the rate at which material is accreted onto a young protostar and the rate at which mass is ejected through bipolar jets or outflows were known, it would be possible to extrapolate the expected outflow rate from this predicted accretion rate.  Unfortunately, no direct measure of the infall/outflow relationship exists, so we must resort to making an educated guess based on the available models. Following \\citet{AndrePPIV} and \\citet{PP}, we assume that realistic jet models give $\\dot{M}_{\\rm jet}/\\dot{M}_{\\rm acc} \\sim 0.1 - 0.3$. With this assumption, we extrapolate that accretion at the rate of $7.4 \\times 10^{-7}$ M$_{\\odot}$ yr$^{-1}$ would be sufficient to power an outflow with an average $\\dot{M}_{\\rm jet} \\sim 0.74-2.2 \\times 10^{-7}$ M$_{\\odot}$ yr$^{-1}$.  In Figure \\ref{fig:mdot}, we have indicated our proposed evolutionary sequence for the 1 M$_{\\odot}$ nonsingular logatropic core from the preprotostellar to the Class I stage.  The preprotostellar core (PPC) stage begins with the core in its near-equilibrium hydrostatic state. The PPC stage ends when the flat-topped central region of the core has collapsed to a singular $r^{-3/2}$ density profile, just at the onset of vigorous accretion.  This marks the end of the collapse phase and the transition to the accretion phase.  The Class 0 stage is characterized by vigorous accretion throughout.  In accordance with the description of \\citet{AndrePPIV}, we mark the transition from the Class 0 to the Class I stage when $M_{\\rm acc} = M_{\\rm env}$.  Once in the proposed Class I stage, the accretion rates of our simulated cores decline rapidly.  This decline in the accretion rate correlates well with the observed absence of strong outflows in Class I cores.  One of the appeals of the logatropic model is its relatively small spread in formation times for stars of a wide range of masses.  The accretion timescales for the formation of 1-5 M$_{\\odot}$ protostars from nonsingular logatropes are on the order of a few million years, for the combinations of $R/r_{0}, P_{s}$, and $T_{c}$ used herein.  These accretion timescales agree well with the observed ages of young star clusters, such as the Orion Nebular Cluster (ONC), wherein the bulk of the stars formed within $\\sim 10^{6}$ yr \\citep{HC2000}.  We anticipate that the accretion timescale will be shorter in regions of higher pressure, as in regions of clustered star formation.  More sophisticated models of collapsing molecular cloud cores will include the effects of rotation and magnetic fields.  The former ensures that most of the material that accretes onto the protostar does so through a circumstellar accretion disk.  The latter, by contrast, is expected to cushion the collapse and perhaps to reduce the infall speeds.  Magnetic fields will also enhance the function of large ``pseudo-disks'' \\citep{GS1} through magnetic focusing of infalling material.  It is important to emphasize that the MP97 models include a measure of magnetic support, in addition to the turbulent support.  The logatrope incorporates a mean magnetic field as a virial parameter, but actual time-dependent simulations employing a full 3-D magnetic field structure are necessary before one can gain a complete picture of logatropic collapse."
    },
    "0201/astro-ph0201526_arXiv.txt": {
        "abstract": "\\noindent We present spin-resolved X-ray data of the neutron star binary Her X-1. We find evidence that the Iron line at 6.4 keV originates from the same location as the blackbody X-ray component. The line width and energy varies over both the spin period and the 35 day precession period. We also find that the correlation between the soft and hard X-ray light curves varies over the 35 day period. ",
        "introduction": "Her X-1 is a stellar binary system consisting of a neutron star and a A/F secondary star. It has been extensively studied at many wavelengths. Its main temporal observational characteristics are: the spin period of the neutron star is $\\sim$1.24 sec, the binary orbital period is 1.7 day and there is a 35 day period seen in X-rays which has been interpreted as due to a warped accretion disc precessing around the neutron star. We have obtained three observations of Her X-1 made using {\\sl XMM-Newton} at 3 different epochs. In this paper we present an analysis of spin resolved data obtained using the EPIC detectors. ",
        "conclusions": "\\subsection{Energy resolved light curves}  Many features of the light curves are naturally explained within the scenario proposed by \\cite{gramsay-c1:scott00}. This model is based on the obscuration of a multi-component X-ray beam by a counter-precessing, tilted, twisted disk.  For simplicity, the X-ray beam is assumed to be decoupled from the disk and is axisymmetric. One of the main features of this model is that it ascribes the variations observed in the pulse profile over the 35~day cycle to occultation from the {\\it inner} part of the disk, whereas most of the previous investigations have assumed an occultation from the {\\it outer} boundary.  The overall situation is summarized in the bottom panel of Figure 8 in \\cite{gramsay-c1:scott00}, while their Figures 10 \\& 11 illustrate the evolution of the pulse profiles predicted during the main-on and short-on respectively. There is a similarity between the multi-peaked hard light curve at $\\Phi_{35}=0.17$ and the model during the progressive occultation of leading and trailing peaks of the hard beam. At $\\Phi_{35} \\sim 0.27$, when the main components are occulted, we only observe the survival of a broad, underlying modulation that is attributed to the magneto-spheric emission. Since this component is emitted from a larger region at some distance from the neutron star, it is naturally expected to have a lower modulation as well as a broad maximum.  The pulse profile close to the short-on is also similar to that presented by \\cite{gramsay-c1:scott00} at $\\Phi_{35}\\sim 0.58$ (see their Figure~11).  In the EPIC data, we can in fact recognize a main peak ``A'' as well a small peak ``B'' (Figure~\\ref{gramsay_c1-fig2}). However, because the notch ``B'' is the hardest feature, spectral considerations suggest that this maximum is associated with the small hard peak and the feature ``A'' with the soft peak discussed by \\cite{gramsay-c1:scott00}. If this is the case, ``B'' is actually due to direct emission from the pencil beam, while ``A'' is the radiation redirected into the fan beam from the antipodal accretion column.  \\subsection{The shift between the soft and hard curves}  Given the complexity of the source, pulse-phase spectroscopy is of paramount importance to separate the different spectral components observed in Her X-1. Using $Einstein$ and $BeppoSax$ data it has been shown that, during the main-on state, the maximum of the thermal component and the power law components are shifted by $\\sim 250^\\circ$. The situation is less consistent as far as the 6.4~keV Fe K line is concerned: \\cite{gramsay-c1:choi94} have shown that its intensity is modulated in phase with the soft emission, suggesting a common origin while \\cite{gramsay-c1:oosterbroek00} have found it correlated with the hard (power law) emission.  The shift in phase between the hard and soft emission can be explained if the latter results from re-processing of hard X-rays in the inner part of the accretion disk. If a non-tilted disk intercepts (and re-processes) a substantial fraction of the hard beam from the neutron star, the expected phase difference between direct and reflected component is $180^\\circ$.  Therefore, the value determined using $Einstein$ and $BeppoSax$ data has been associated with the disk having a tilt angle. If the tilt of the disk changes with the phase along the 35~day cycle (as predicted by the precessing disk models, see \\cite{gramsay-c1:gerend76}) the shift in phase should therefore vary with $\\Phi_{35}$. However, both Einstein and Sax data were obtained at the same $\\Phi_{35}$.  The phase shift derived from {\\sl XMM-Newton} data main-on state data are considerably different from previous observations made in the main-on state and continues to change during the other two observations. This suggests that we are observing, a {\\it substantial and continuous variation in the tilt of the disk}, which is what we would expect from a system which had a precessing accretion disc. It should be noted that the interpretation of the phase shift observed at the short-on may be affected by a systematic error, depending on whether during the observation the soft peak ``A'' is higher than the small hard peak ``B'' or vice-versa.  \\subsection{The Fe 6.4 keV line variation} \\label{line}  At $\\Phi_{35}=0.26$, there is little evidence for a significant variation in the Fe line parameters, while at $\\Phi_{35}=0.60$, there is some evidence that the variation of the equivalent width of the Fe line is in anti-phase with the (2--4)~keV intensity curve and follows the general shape of the (0.3--0.7)~keV intensity curve.  At $\\Phi_{35}=0.17$, we find that the soft flux below 0.7~keV, the line normalization, the line width and the equivalent width all exhibit a common minimum at $0.7<\\phi_{spin}<0.9$, which, in turn, is shifted with respect to that of the hard emission. This supports the idea that the 6.4~keV Fe line originates from fluorescence from the relatively cold matter of the illuminated spot where the soft emission is reprocessed.  We have also found evidence for a variation in the Fe line broadening over the 35~day period. In addition, the energy of the Fe line suggests that the Fe line emission originates from low ionisation species (Fe XIV or less) in the low and short-on state observations, whereas in the main-on the observed Fe K centroid energies ($6.52 \\pm 0.03$~keV for MOS and $6.50 \\pm 0.02$~keV for PN) correspond to Fe XX-Fe XXI.  Taking this into account, the true centroids deviate by $\\sim7 \\sigma$ from the 6.40~keV neutral value. This suggests two possible explanations for both the line broadening and the centroid displacement: 1) an array of Fe K fluorescence lines exists for a variety of charge states of Fe (anything from Fe I-Fe XIII to Fe XXIII); 2) Comptonization from a hot corona with a significant optical depth for a narrower range of charge states centered around Fe XX.  The Fe line broadening may also be explained in terms of Keplerian motion, if the inner disk (or some inner region) comes into view during the main-on state. If this is the case, at $\\Phi_{35}=0.17$ the Keplerian velocity will be $\\sim 13000$ km/sec. This, in turn, corresponds to a radial distance of $\\sim 4 \\times 10^8$~cm (for a neutron star of 1.4\\sun), which is close to the magneto-spheric radius for a magnetic field of $\\sim 10^{12}$~G."
    },
    "0201/astro-ph0201183_arXiv.txt": {
        "abstract": "The prescient remark by Baade and Zwicky that supernovae beget neutron stars did little to prepare us for the remarkable variety of observational manifestations such objects display. Indeed, during the first thirty years of the empirical study of neutron stars, only a handful were found to be associated with the remnants of exploded stars. But recent X-ray and radio observations have gone a long way toward justifying the theoretical link between supernovae and neutron stars, and have revealed the wide range of properties with which newborn compact remnants are endowed. We review here our current state of knowledge regarding neutron star-supernova remnant associations, pointing out the pitfalls and the promise which such links hold. We discuss work on the ranges of neutron star velocities, initial spin periods, and magnetic field strengths, as well as on the prevalence of pulsar wind nebulae. The slots in neutron star demography held by AXPs, SGRs, radio-quiet neutron stars, and other denizens of the zoo are considered. We also present an attempt at a comprehensive census of neutron star-remnant associations and discuss the selection effects militating against finding more such relationships. We conclude that there is no pressing need to invoke large black hole or silent neutron star populations, and that the years ahead hold great promise for producing a more complete understanding of neutron star birth parameters and their subsequent evolution. ",
        "introduction": "The prediction by Baade \\& Zwicky (1934) that neutron stars would form in supernova explosions ranks among the most prophetic theoretical speculations in the history of astrophysics. An important aspect of the matter these two illustrious astronomers did not address was how such beasts could be observed.  Certainly the properties of the rotation-powered pulsars discovered by Jocelyn Bell were unanticipated before 1967; the thermal X-ray emission detected from the surfaces of a handful of nearby young pulsars over a quarter of a century later (e.g.\\ Finley, \\\"Ogelman \\& Kiziloglu 1992) was the first properly predicted neutron-star observational characteristic.  But the unanticipated reared its head in the interim, and the detection of soft gamma-ray bursts from the supernova remnant N49 in the Large Magellanic Cloud in 1979 (Mazets et al.\\ 1979; Cline et al.\\ 1982), as well as the discovery, from a source in the remnant CTB~109, of slow X-ray pulsations that could be explained neither in terms of rotation-power nor in terms of any conventional accretion mechanism (Fahlman \\& Gregory 1981), suggested that the words ``unexpected zoo'' may be applicable to the observed young neutron star population.   This diversity of appearance continues to grow: X-ray point sources with properties quite different from those described above include the variable point source in RCW~103 (Tuohy \\& Garmire 1980; Gotthelf, Petre \\& Vasisht 1999; Garmire et al.\\ 2000), the slow and X-ray variable pulsar AX J1845$-$0258 in G29.6+0.1 (Vasisht et al.\\ 2000), and the point source in Cas~A, made famous by the spectacular first light {\\em Chandra X-ray Observatory}\\ image (Tananbaum 1999).  It is perhaps ironic that in spite of the many possible observational manifestations of young neutron stars, one of the greatest problems in the field has been that most Galactic supernova remnants appear empty. Observational selection effects must certainly play a role, but ever-lingering is the fact that an unknown but possibly substantial fraction of massive stars undergoing core collapse might produce black holes rather than neutron stars.  The size of this fraction is observationally poorly constrained.  Certainly the difficulty that supernova modelers encounter in producing shock waves that can expel the outer layers of a star following the bounce in core collapse does little to dispel the concern that the black hole fraction may be large (Burrows, Hayes \\& Fryxell 1995; Liebend\\\"order et al.\\ 2001; Janka 2001).  Our ignorance of the neutron star equation of state and hence the maximum stable neutron star mass adds to this confusion (Baym \\& Pethick 1979).  A radio pulsar is thought to be born between every 60 and 330~yr in the Galaxy (Lyne et al.\\ 1998).  This is in rough agreement with the best-estimate core-collapse Galactic supernova rate of one every $47 \\pm 12$~yr (Tammann, L\\\"offler \\& Schr\\\"oder 1994), although a significant paucity of pulsars is suggested, and the uncertainties on both quantities are large enough to leave plenty of room for suspicion.  In this review, we discuss first the associations between SNRs and classical rotation-powered neutron stars. (Henceforth, we refer to the latter simply as ``pulsars''; non-rotation-powered pulsars, in particular anomalous X-ray pulsars, will be specifically identified.) Our emphasis is less on a critique of all proposed associations (since significant reviews of this nature already exist --- see e.g.\\ Kaspi 1998, 2000; Manchester 1998) and more on what fundamental astrophysical insights such associations provide.  Indeed, pulsar/supernova remnant associations have been heralded as, in principle, constraining the neutron star initial spin period, velocity and magnetic field distributions, all important for the physics of core collapse.  In addition, dating pulsars is valuable in constraining the equation of state by testing cooling models (e.g.\\ Umeda et al.\\ 1994) --- associations with supernova remnants can in principle offer independent age determinations.  We subsequently consider associations between supernova remnants and non-rotation-powered neutron stars as well as compact remnants of an uncertain nature, not merely for completeness, but as a precursor to the final section, in which we present an attempt at a synthesis of all associations.  The goal here is to address the larger issue of core-collapse-remnant demographics.  Similar attempts have been made in the past (e.g.\\ Helfand \\& Becker 1984; Helfand 1998), although the new observational capabilities in the high energy regime represented by {\\em Chandra}\\ and {\\em XMM-Newton}\\ have led, in the last two years, to a dramatic acceleration in our identification of the compact remnants of supernova explosions. Indeed, we conclude here that current data are consistent with the notion that the majority of core-collapse events in the Galaxy {\\it do} leave neutron star remnants. ",
        "conclusions": "As with any maturing astronomical field (or in life, for that matter), the simple pictures of youth grow with time more complicated and more frustrating --- or richer, depending on one's personal gestalt. The triumphalist picture of science as proceeding from an inspired insight (Baade and Zwicky 1934) to a more detailed model (Pacini 1967) to an experimental confirmation (Hewish et al.\\ 1968) has always been a trifle simplistic. Supernovae do not {\\it always} ``represent the transition from an ordinary star into a neutron star'', and the ones that do evidently provide us with a host of possible manifestations. We have presented here a snapshot of our current state of knowledge concerning the genesis of neutron stars. The fact the we cite no fewer than 39 references from within the past twenty-four months suggests the subject is a vibrant one which is benefiting from significant improvements in observational sensitivity, primarily in the X-ray and radio bands. It took thirty years to find the first ten secure pulsar-remnant associations and less than three years to find the next five. The large number of candidates discussed above assures continuing rapid progress. The longstanding embarrassment of all those empty supernova remnants appears to be significantly ameliorated and is replaced by the challenge of overcoming the selection effects which hamper discoveries, and understanding the many ways in which young neutron stars manifest their presence.  Such understanding might even lead to quantitative constraints on the conditions present in core-collapse explosions and the resulting nucleosynthetic yields which determine the chemical composition of the Galaxy."
    },
    "0201/astro-ph0201460_arXiv.txt": {
        "abstract": "The AGN-galaxy cross-correlation function of radio-quiet AGN and radio galaxies has been measured with the Panoramic Deep Fields. Colour selection criteria and photometric redshifts have been used to significantly increase the signal-to-noise of the angular cross-correlation function. Radio-quiet AGN environments are comparable to the environments of early-type galaxies at low redshift. The radio galaxy-galaxy spatial cross-correlation function is very strong though large variations are observed from field-to-field. These variations appear to be caused by large-scale-structure on scales comparable to the $5^\\circ \\times 5^\\circ$ field-of-view. The distribution and spatial cross-correlation function of radio galaxies and clusters in the Panoramic Deep Fields is consistent with these objects tracing the same structures at $z<0.7$. No evidence is found for evolution of the AGN-galaxy spatial cross-correlation function across the redshift range observed. ",
        "introduction": "The environments of radio-quiet AGN and radio galaxies have been measured using the Panoramic Deep Fields, a $UB_JRI$ imaging survey of two $5^\\circ \\times 5^\\circ$ fields with a depth of $B_J\\sim 23.5$ and $R\\sim 22.5$. The wide field-of-view and depth was obtained by coadding SuperCOSMOS scans of UK Schmidt photographic plates. The resulting galaxy catalogues contain more than $2\\times 10^5$ galaxies per field and thus provide the large sample size required for accurate estimates of AGN-galaxy clustering. As the galaxy catalogues contain more than 1000 objects with spectroscopic redshifts from the NED database, the polynomial fitting technique of Connolly {\\it et al.} (1995) was used to estimate photometric redshifts for all galaxies with $B_J$ and $R$-band detections. The B1950 coordinates of the field centres are $00~53~-28~03$ (SGP field) and $10~40~+00~00$ (F855 field). Access to the galaxy catalogue and cutout images are available via {\\tt http://astro.ph.unimelb.edu.au/data/}. ",
        "conclusions": "The Panoramic Deep Fields have been used to measure the AGN-galaxy cross-correlation function for radio-quiet AGN and radio galaxies. By applying colour selection and photometric redshift criteria, it has been possible to significantly increase the signal-to-noise of the angular cross-correlation function. The clustering of red (early-type) galaxies around $UB_JR$ selected AGN is comparable to the clustering of early-type galaxies at low redshift. The clustering of galaxies around radio galaxies is strong though large variations in the clustering strength are observed between the two Panoramic Deep Fields. The distribution and spatial cross-correlation function of radio galaxies and galaxy clusters indicate that these objects trace the same large-scale-structures."
    },
    "0201/hep-ph0201155_arXiv.txt": {
        "abstract": "We discuss various issues related to stabilized embedded strings in a thermal background. In particular, we demonstrate that such strings will generically become superconducting at moderately low temperatures, thus enhancing their stability. We then present a new class of defects - drum vortons - which arise when a small symmetry breaking term is added to the potential. We display these points within the context of the $O(4)$ sigma model, relevant for hadrodynamics below the QCD scale. This model admits `embedded defects' (topological defect configurations of a simpler - in this case $O(2)$ symmetric - model obtained by imposing an embedding constraint) that are unstable in the full model at zero temperature, but that can be stabilised (by electromagnetic coupling to photons) in a thermal gas at moderately high termperatures. It is shown here that below the embedded defect stabilisation threshold, there will still be stabilized cosmic string defects. However, they will not be of the symmetric embedded vortex type, but of an `asymmetric' vortex type, and are automatically superconducting. In the presence of weak symmetry breaking terms, such as arise naturally when using the $O(4)$ model for hadrodynamics, the strings become the boundary of a new kind of cosmic sigma membrane, with tension given by the pion mass. The string current would then make it possible for a loop to attain a (classically) stable equilibrium state that differs from an ``ordinary'' vorton state by the presence of a sigma membrane stretched across it in a drum like configuration. Such defects will however be entirely destabilised if the symmetry breaking is too strong, as is found to be the case -- due to the rather large value of the pion mass -- in the hadronic application of the O(4) sigma model. ",
        "introduction": "\\label{sec:1}  The purpose of this work is to follow up the work of Nagasawa and Brandenberger~\\cite{Nagasawa:1999iv} who considered the possibility of thermal stabilisation, via electromagnetic coupling, of vortex defects, i.e. cosmic strings, in a Sigma model characterised by O(4) symmetry with a set of degenerate vacuum states having the topology of a 3-sphere.  Since the homotopy structure of the 3 sphere is trivial, such a model does not have stationary vacuum defects of a topologically stable kind. However this model (involving charged and neutral pion fields as well as the sigma field) contains a subset of solutions that is identifiable as the complete set of solutions of an ``embedded'' model (involving just the neutral pion and the sigma field) characterised by O(2) symmetry. This embedded model has a set of degenerate vacuum solutions having the topology of a circle, and therefore admits stationary vacuum vortex defects of a topologically stable kind, which were called {\\it pion strings} in the initial paper \\cite{Zhang:1997is} on this subject \\footnote{In this paper we will restrict our attention to the classical Sigma model and not touch on the rich variety of defects which can exist when the quantum nature of QCD (in particular at high baryon density) is taken into account (see e.g. \\cite{Halperin:1998gx,Forbes:2000et,Son:2000fh,Son:2001xd,Kaplan:2001hh,Forbes:2001gj,Balachandran:2001qn} for discussions of such defects).}. These stationary topological defect configurations of the embedded O(2) model constitute what are known \\cite{Vachaspati:dz,Achucarro:1992hs,Barriola:fy,Achucarro:1999it} as embedded defects within the framework of the full model, but as their energy is not minimised in the broader framework of the full O(4) model they will not be stable in this more general context.  The point made by Nagasawa and Brandenberger \\cite{Nagasawa:1999iv} was that the background reference states that are relevant in cosmological contexts are commonly not vacuum states but thermal equilibrium states, for which topological defects of the embedded O(2) model can be stable as vortex defects  of the full model. The possibility of creating such vortex defects, i.e. cosmic strings, arises from breaking of the  O(4) symmetry by thermal effects mediated by electromagnetic coupling. Such stabilisation of an embedded defect (i.e. of a topological defect of the embedded O(2) symmetric model) does however require that the product of the relevant electric coupling constant $e$ and the temperature $\\Th$ should be sufficiently large.  The first thing we wish to point out here is that topologically stabilised vortex defects of thermal (not vacuum) equilibrium states will exist for any non zero value of the product $e\\Th$, even if it is very small (as long as the temperature is higher than the temperature of recombination, below which the thermal analysis used in this paper breaks down). For large values of $e\\Th$ these topological defects include the embedded defects referred to above. However for smaller values of $e\\Th$ the topological defects are not configurations of the embedded model, but are of a mathematically less trivial kind with the important property that (unlike their embedded counterparts at higher temperature) they are automatically ``superconducting'' in the sense of Witten~\\cite{Witten:eb}. As first observed by Davis and Shellard~\\cite{Davis:ij}, such a conductivity property allows cosmic string loops to form vortons, i.e. centrifugally supported equilibrium states, which under a wide range of conditions will actually be stable~\\cite{Carter:wu}.   The foregoing considerations are based on the supposition that the underlying field model has non thermal vacuum states characterised by strict O(4) symmetry, with respect to which the pions are identifiable as Goldstone bosons which as such must have zero mass. However for a more realistic description, allowing for a finite pion mass that is actually observed, the Lagrangian of the model has to be augmented by the inclusion of a small intrinsic O(4) symmetry breaking term. This removes the degeneracy of the vacuum, as well as of the thermal equilibrium states, so there is no longer any possibility of forming a topologically stable defect, whether of the vacuum or of a thermal equilibrium state at finite temperature.  There is however the possibility at finite temperature of setting up a stationary state of a more interesting kind. One of the purposes of this article is to consider the construction in such a context of a more general kind of (dynamically but not topologically) stable equilibrium configuration  that may be described as ``drum vorton'' (or ``frisbee'')  consisting of a vorton like loop forming the boundary of a drum type membrane.  It is shown that the existence of such stabilised defects is only possible if the symmetry breaking term is sufficiently small. This condition may be satisfied in other applications, but it is found that it does not hold in the case when the O(4) sigma model is applied in the hadrodynamic context for which it was originally designed. The failure of the stabilisation mechanism in this particular case is attributable to the rather large value of the (destabilising) pion mass $m_\\pi$ in conjunction with the rather small value $e^2 \\simeq 1/137$ of the (stabilising) electromagnetic coupling constant. ",
        "conclusions": "\\label{sec:14} In this paper we have studied the stabilization mechanism for embedded defects \\cite{Nagasawa:1999iv}, with particular emphasis on the application to the classical bosonic $O(4)$ sigma model of hadrodynamics \\footnote{After completion of this manuscript a preprint \\cite{Ward:2002ci} appeared which discusses the stabilization of certain unstable strings and textures by the cosmological expansion.}.  We have seen that below the stabilization threshold for an embedded defect of the traditional kind (with symmetric core) there will still be stablized cosmic string defects, but of asymmetric vortex type. These defects will automatically be superconducting, and this provides them with an extra stabilization mechanism. These superconducting string defects are stable above a threshold temperature $\\Th_d$ set by the strength of the explicit symmetry breaking term in the potential, i.e. by the pion mass in the case of hadrodynamics. In the absence of explicit symmetry breaking the defects remain stable until the temperature of recombination, at which point our thermal analysis breaks down.  In the case of explicit symmetry breaking, the superconducting vortices become boundaries of a new type of membrane-like defects which we call {\\it drum vortons}, across which the change in the phase of the string order parameter is localized, and whose tension is given by the symmetry breaking mass, the pion mass in the case of hadrodynamics.  We have seen that drum vortons can be stabilized by rotation.  In the case of hadrodynamics, the pion mass is too large for the superconducing vortices and drum vortons studied here to be stable. This is due to the large value of the pion mass relative to the QCD symmetry breaking scale, and due to the large value of the self coupling constant $\\lambda$ relative to the small value of the gauge coupling $e^2$. However, in many Grand Unified Models, we expect $\\lambda$ to be small, and the explicit symmetry breaking terms to be absent. In this case, the embedded strings with asymmetric core studied in this paper and their drum vortons would be stable.  Thus, we have identified a new class of defects which could be of great cosmological importance in the early Universe. They could be used for baryogenesis (see e.g. \\cite{Davis:1996zg}) or for the generation of primordial magnetic fields (see e.g. \\cite{Brandenberger:1998ew}). There are also severe cosmological constraints on models which admit such defects, a topic which we will come back to in a subsequent publication \\cite{BCD2}.   \\centerline{Acknowledgements}  The authors wish to thank J. Blanco-Pillado and A. Zhitnitsky for discussions. BC and ACD thank Brown University and RB and ACD thank UBC for hospitality whilst this work was in progress. This work was supported in part by the ESF COSLAB programme and by a Royal Society-CNRS exchange grant (BC,ACD), by PPARC (ACD), by the US Department of Energy under Contract DE-FG0291ER40688, Task A (RB), and by an Accord between CNRS and Brown University (BC,RB). We are grateful to Herb Fried for securing this Accord."
    },
    "0201/astro-ph0201240_arXiv.txt": {
        "abstract": "{ A few Seyfert 1s have a H$\\beta$ profile with a red wing usually called the ``red shelf\". The most popular interpretation of this feature is that it is due to broad redshifted lines of H$\\beta$ and [O~III]$\\lambda\\lambda$4959,5007; we have observed two Seyfert 1s displaying a ``red shelf\" and showed that in these two objects the main contributor is most probably the He~I $\\lambda\\lambda$4922,5016 lines having the velocity and width of the broad H$\\beta$ component. There is no evidence for the presence of a broad redshifted component of H$\\beta$ or [O~III] in any of these two objects. ",
        "introduction": "A few Seyfert 1s have a very complex H$\\beta$ profile with a strong red wing extending underneath the [O~III]$\\lambda\\lambda$4959,5007 lines. The excess emission in the red wing is referred to as the ``shelf\" feature or the ``red shelf\" (Meyers \\& Peterson \\cite{meyers}). In most Seyfert 1s showing such a feature, it appears to be made of two components: a broad red wing to H$\\beta$, and a broad wing on the long wavelength side of the [O~III]$\\lambda$5007 line (van Groningen \\& de Bruyn \\cite{groningen}).  Several interpretations have been proposed to explain the ``red shelf\": it could be due to H$\\beta$ or to the presence of other broad emission lines such as [O~III], Si~II $\\lambda$5056, He~I $\\lambda$5016 (Meyers \\& Peterson \\cite{meyers}; van Groningen \\& de Bruyn \\cite{groningen}; Kollatschny et al. \\cite{kollatschny01}) or Fe~II (Korista \\cite{korista92}).  Meyers \\& Peterson (\\cite{meyers}), Crenshaw \\& Peterson (\\cite{crenshaw1986}) and Stirpe et al. (\\cite{stirpe89}) have argued that broad [O III] lines are most probably the main contributor to the observed $\\lambda$5007 red wing. Van Groningen \\& de Bruyn (\\cite{groningen}) have found the same red wing in [O~III]$\\lambda$4363 in several objects, confirming that they are indeed due to [O~III] emission.  To investigate the nature of the ``red shelf\", we have made spectroscopic observations of two Seyfert 1s: RXS~J01177+3637 and HS~0328+05. ",
        "conclusions": "The detailed study of the spectrum of two Seyfert 1s showing a ``red shelf\" shows that it is mainly due, in addition to the Fe~II multiplet 42, to the presence of relatively strong He~I $\\lambda$4922 and $\\lambda$5016 broad lines. In HS~0328+05, there is also a non negligible contribution to the H$\\beta$ red wing of a number of narrow emission lines.  There is no evidence for the presence of a broad redshifted component in H$\\beta$ or [O~III] in any of these two objects."
    },
    "0201/astro-ph0201289_arXiv.txt": {
        "abstract": "We study the phase-space structure of a dark-matter halo formed in a high resolution simulation of a $\\Lambda$CDM cosmology.  Our goal is to quantify how much substructure is left over from the inhomogeneous growth of the halo, and how it may affect the signal in experiments aimed at detecting the dark matter particles directly.  If we focus on the equivalent of ``Solar vicinity'', we find that the dark-matter is smoothly distributed in space. The probability of detecting particles bound within dense lumps of individual mass less than $10^7 M_\\odot h^{-1}$ is small, less than $10^{-2}$. The velocity ellipsoid in the Solar neighbourhood deviates only slightly from a multivariate Gaussian, and can be thought of as a superposition of thousands of kinematically cold streams. The motions of the most energetic particles are, however, strongly clumped and highly anisotropic. We conclude that experiments may safely assume a smooth multivariate Gaussian distribution to represent the kinematics of dark-matter particles in the Solar neighbourhood.  Experiments sensitive to the direction of motion of the incident particles could exploit the expected anisotropy to learn about the recent merging history of our Galaxy. ",
        "introduction": "One of the most fundamental open questions in cosmology and particle physics today is what is the nature of dark-matter. The first indications of its existence came in the 1930s, with the measurements of the velocities of galaxies in clusters. The cluster mass required to gravitationally bind the galaxies was found to be roughly an order of magnitude larger than the sum of the luminous masses of the individual galaxies \\cite{zwicky,smith}.  In the 1970s, observations of the rotation curves of spiral galaxies ($V_c(r) = \\sqrt{GM(r)/r}$) showed that these were flat or even rising at distances far beyond their stellar and gaseous components \\cite{rubin70,faber79,rubin80}. These discoveries led to the conclusion that a large fraction (more than 90\\%) of the mass in the Universe is dark. It is now widely believed that this mass is most likely in the form of yet to be discovered nonbaryonic elementary particles.  Being the dominant mass component of galaxies and of large-scale structures in the Universe, dark-matter has necessarily become a key ingredient in theories of structure formation in the Universe. The most successful of these theories is the hierarchical paradigm \\cite{peebles74}. In the current (and observationally most favoured) version of this model, the nonbaryonic elementary particles are known as ``cold dark-matter\" (CDM) \\cite{peebles82}.  The term ``cold\" derives from the fact that the dark-matter particles had non-relativistic motions at the time of matter--radiation equality. Their abundance was set when the interaction rate became too small for the particles to be in thermal equilibrium in the expanding Universe.  The first objects to form in a CDM Universe are small galaxies, which then merge and give rise to the larger scale structures we observe today. Thus structure formation occurs in a ``bottom-up\" fashion {\\cite{blumenthal84,frenk83}.  This hierarchical paradigm has allowed astronomers to make very definite predictions for the properties of galaxies today and about their evolution from high redshift. Direct comparisons to observations have shown that this model is quite successful in reproducing both the local and the distant Universe.  The crucial test of this paradigm undoubtedly consists in the determination of the nature of dark-matter through direct detection experiments. Among the most promising candidates from the particle physics perspective are axions and neutralinos.  Axions have been introduced to solve the strong-CP (Charge conjugation and Parity) violations \\cite{PC}. They can be detected through their conversion to photons in the presence of a strong magnetic field (e.g. \\cite{ax1,ax2}). Neutralinos are the lightest supersymmetric particles, and may be considered as a particular form of weakly interacting massive particles (WIMPs). The most important direct detection process of neutralinos is through elastic scattering on nuclei. The idea is to determine the count rate over recoil energy above a given (detector) background level. The experimental situation has been improving rapidly over the past years, with large-scale collaborations such as DAMA, Edelweiss and CDMS ~\\cite{DAMA00,CDMS00,EDELWEISS} starting to probe interesting regions of parameter space (for an extensive discussion see \\cite{Berg}). The main problem currently lies in the high level of background noise, either from ambient radioactivity or cosmic-ray induced activity. Information on the direction of the recoils could potentially also be useful and yield a large improvement in sensitivity \\cite{Spergel}.  In all these experiments, the count rate strongly depends on the velocity distribution of the incident particles, and a modulation effect due to the orbital motion of the Earth around the Sun is expected \\cite{Drukier}. In most cases, an isotropic Maxwellian distribution has been assumed (e.g. ~\\cite{Freese88,Bernabei98}), although there are other examples in the recent literature, discussing multivariate Gaussian distributions \\cite{Ullio01,Green01,wyn01}. Attempts at understanding the effect of substructure in the velocity distribution of dark-matter particles have also been made \\cite{sikivie98,stiff01,hogan01}. This substructure would have its origin in the different merger and accretion events the Galaxy should have experienced over its lifetime \\cite{hw99}.  The progressive build up of dark halos through mergers and accretion of smaller subunits implies that the latter will leave substructure in the phase-space of the final object. This is because the phase-space volume of the final object is much larger than that initially available for each one of the objects independently. For example, for a small satellite galaxy the initial phase-space volume occupied by its particles is proportional to $(R^{\\rm sat} V_c^{\\rm sat})^3 $, where $R^{\\rm sat}$ is the size of the satellite, and $V_c^{\\rm sat}$ its circular velocity. The volume available to the satellite particles after the merging is determined by their orbit, and is a factor $(R^{\\rm gal}/R^{\\rm sat})^3 \\times (V_c^{\\rm gal}/V_c^{\\rm sat})^3$ larger, where $R^{\\rm gal}$ is the size of the final object and $V_c^{\\rm gal}$ its circular velocity. Note that even in the case of a major merger, where the mass is doubled, the phase-space volume available is already 4 times larger \\footnote{This is because $M \\propto R^3 \\propto V^3$, and thus if $M_f = 2 M_i$ then $ R_f^3 V_f^3 \\propto 4 R_i^3 V_i^3$}. The key question is whether this substructure will be directly or indirectly observable. For example, if there is a bound satellite going through the Solar neighbourhood at the present day, it will dominate the flux of dark-matter particles on Earth. The energy spectrum of these particles will be strongly peaked around the orbital energy of the clump, perhaps giving a signal similar to a delta function.  As we shall show in Sec.\\ref{sec:boundhalos} the fraction of mass in satellites which could have survived the tidal field of the Galaxy by the present day is less than $10^{-2}$ of total mass of the Galaxy, implying that such a scenario is relatively unlikely.  More realistic is to assume that the satellite halos that contribute with mass to the Solar neighbourhood will be completely disrupted. The particles freed from such satellites will tend to follow the initial orbit of their progenitor, and eventually will fill a volume comparable to the size of the orbit.  Because of the conservation of phase-space density (Liouville's theorem), this implies that locally they should have very similar velocities \\footnote{If initially $\\Delta_x \\Delta_v$ is the phase-space volume occupied by the satellite, and if $\\Delta_x'$ is its final volume, then $\\Delta_v' = \\Delta_v \\times \\Delta_x'/\\Delta_x$, where as discussed above, $\\Delta_x'$ is the volume given by the orbit, and is much larger than the original volume of the satellite.}.  Thus one may expect to see streams of particles going through the Solar neighbourhood, which had their origin in the different merging events.  Such streams have already been observed in the motions of nearby halo stars and in the outer regions of the Galactic halo \\cite{helmi99,ibata94}. Streams manifest themselves as peaks in the velocity distribution function. Clearly it is important to determine for the dark-matter particles in the vicinity of the Sun whether this distribution function will be dominated by a few of these peaks, or whether their number is so large, that it will be close to Gaussian.  The best way to understand the expected properties of the Galactic halo in the Solar neighbourhood is through high-resolution simulations starting from appropriate cosmological initial conditions.  Analytic modelling can provide insights into the processes that drive the build-up of structure such as phase-mixing, or tidal stripping. Nevertheless it needs to be complemented by cosmological simulations, that provide the mass spectrum of the accreted halos, their orbital parameters, their characteristic merging times, and the detailed mixing of the material they deposit. The highly non-linear character of the hierarchical build up of a galaxy like the Milky Way, forces us to resort to numerical simulations to make realistic predictions for its properties. Very high resolution simulations are required to be able to resolve the substructures leftover from merging events, since their density contrasts are expected to fade rather quickly with time (as $t^3$ for sufficiently long timescales \\cite{hw99}).  The main goal of the present paper is to understand the phase-space structure of a dark-matter halo. We wish to quantify the expected amount of substructure and understand its effect on direct dark-matter detection experiments.  Particular emphasis will be put on determining the properties of the dark-matter distribution in the Solar neighbourhood: its mass growth history, the spatial distribution and the kinematics of particles in this region of the Galactic halo.  We address these issues by scaling down a high-resolution simulation of the formation of a cluster of galaxies in a $\\Lambda$CDM cosmology to a galactic size halo \\cite{vrs01a}. ",
        "conclusions": "We analysed a high resolution simulation of the formation of a cluster in a $\\Lambda$CDM cosmology. By scaling it down to a galaxy size halo (by the ratio of the maximum circular velocities) we were able to make predictions for the expected dark-matter distribution near the Sun.  Our results indicate that direct detection experiments may quite safely assume that the distribution of dark-matter particles in the Solar neighbourhood is well represented by a multivariate Gaussian. We find that none of the streams present in any of the volumes at the Sun's distance from the Galactic centre dominate their local distribution. The mean density of an individual stream is typically 0.3\\% that of the local dark-matter distribution (deduced from the number of particles in the rather dense stream shown in Figure~\\ref{fig:DM_4kpc}).  These small values are due to the fact that most of the streams in the inner galaxy come from a few massive halos that merged at high redshift to build up the object we see today. These large halos mix extremely quickly and therefore give rise to low density structures. Strong density enhancements such as those predicted in \\cite{sikivie98, sikivie01} are extremely unlikely in the inner Galaxy.  Our simulation also shows that we should not expect to find dense, recently formed streams near the Sun, since the last accretion contributing matter to the Solar neighbourhood typically took place about 1 Gyr ago, and provides only $\\sim 10^{-4}$ of the total mass in this region. Moreover, we find that fewer than 1\\% of the local dark-matter particles could be part of small dense subhalos which have survived intact within the larger halo of the Milky Way. It is therefore unlikely that an individual halo with these characteristics will dominate the signal in direct detection experiments.  Direct detection experiments which are sensitive to the direction of motion of the fastest moving dark-matter particles may discover a direct indication of the hierarchical growth of our Galaxy's halo. The expected signal for these fastest moving particles is highly anisotropic, and could be eventually be used, not just to determine the nature of the dark-matter, but also to recover, at least partially, the recent merging history of the Milky~Way.  We thank Ben Moore and Leo Stodolsky for useful discussions which triggered many of the ideas discussed here; Uros Seljak and Anne Green for helping us improve the manuscript. AH wishes to thank Joke for being a continuous source of inspiration.  \\newpage"
    },
    "0201/gr-qc0201031_arXiv.txt": {
        "abstract": "The quantum fluctuation of the relative location of two $(n-1)$-dimensional de~Sitter branes (i.e., of $n$ spacetime dimensions) embedded in the $(n+1)$-dimensional anti-de Sitter bulk, which we shall call the quantum radion, is investigated at the linear perturbation level. The quantization of the radion is done by deriving the effective action of the radion. Assuming the positive tension brane is our universe, the effect of the quantum radion is evaluated by using the effective Einstein equations on the brane in which the radion contributes to the effective energy momentum tensor at the linear order of the radion amplitude. Specifically, the rms effective energy density arising from the quantum radion is compared with the background energy density. It is found out that this ratio remains small for reasonable values of the parameters of the model even without introducing a stabilizing mechanism for radion, although the radion itself has a negative mass squared and is unstable. The reason behind this phenomenon is also 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 either of two boundary 3-branes with positive and negative tensions in the 5-dimensional anti-de Sitter space (AdS) \\cite{RS1,RS2}. One of the attractive features of the Randall-Sundrum (RS) scenario is that the gravity on the brane is confined within a short distance from the brane even for an infinitely large extra-dimension\\cite{RS2,GT}. This applies to the positive tension brane and it is because the AdS bulk on both sides of the positive tension brane shrinks exponentially as one goes away from the brane.  Since the RS scenario gives an exciting, new picture of our universe, it is clearly important to study various aspects of this scenario to test it or to give constraints on its parameters. One good example is an analysis done by Garriga and Tanaka\\cite{GS}, in which they have shown that the radion in the original RS scenario acts as a Brans-Dicke scalar on the branes at the linear perturbation order and the effective gravity is that of a Brans-Dicke theory with positive and negative Brans-Dicke parameters on the positive and negative tension branes, respectively, where the values of the Brans-Dicke parameters are determined by the distance between the two branes. In a previous paper, we have shown that essentially the same situation arises in the case of two de Sitter (dS) branes embedded in the AdS bulk. However, we have also shown that the radion effectively has a negative mass squared with its absolute value proportional to the curvature of the dS brane, hence is unstable if it can fluctuate by itself without the matter energy momentum tensor.  The phenomenon we shall study in this paper is the quantum fluctuation of this mode, called the quantum radion. To make clear what we mean by the quantum radion, let us describe the mode in more detail. Our brane universe may be displaced from the 0th order trajectory of a homogeneous and isotropic brane. By appropriately fixing the coordinate gauge, the displacement perpendicular to the brane can be described by a scalar function on the brane. Assuming there are two branes that are fixed points of the $Z_2$ symmetry, it can be shown that only the relative displacement of the two branes is physical, which we call the radion. There are two distinctively different kinds of displacement of a brane: the ``bend'' and ``fluctuation''. The bend is a type of displacement due to inhomogeneities of the matter energy-momentum tensor on the brane. The trace of the energy-momentum acts as an additional tension of the brane, and the brane must ``bend'' accordingly. The relative bend is described by the mode of radion that couples with the source on the branes, and it can be written as a functional of the energy-momentum tensor\\cite{GT,GS}. In the context of the quasi-localized gravity discussed by Gregory, Rubakov \\& Sibiryakov\\cite{GRS}, the role of this type of radion has been extensively studied\\cite{CEH1,DGP,CEH2,PRZ,KMPP}. On the other hand, the ``fluctuation'' is a type of displacement that is purely geometrical, which obeys a free wave equation without source. The relative displacement of this kind is the mode of radion which we shall discuss in this paper. This mode of radion was studied first by Charmousis, Gregory and Rubakov for the RS branes whose effective radion mass is zero by solving the field equations for the RS branes \\cite{CGR} and by Chacko and Fox\\cite{CF} for the dS and AdS branes whose effective radion mass squared are negative and positive, respectively.  The fact that the radion mass squared is negative (or zero) suggests the (marginal) instability of the two brane system. In fact, in the case of the original RS flat two brane model\\cite{RS1}, the negativity of the Brans-Dicke parameter on the negative tension brane\\cite{GT} can be regarded as a result of this marginal instability. To recover the stability, Goldberger \\& Wise introduced a bulk scalar field\\cite{GW} that couples to the branes in such a way that the distance between the two branes is stabilized. However, it should be also noted that the effective Brans-Dicke parameter on the positive tension brane is positive and it can be large enough to be consistent with experiments for the separation of the branes larger than the AdS curvature radius at least at the linear perturbation order\\cite{GT}. Hence a stabilization mechanism may be unnecessary if we live on the positive tension brane. We therefore do not introduce a stabilization mechanism.  We shall work on the system that consists of an $(n+1)$-dimensional AdS bulk spacetime bounded by two branes of constant curvature that are fixed points of the $Z_2$-symmetry. The zero-curvature branes correspond to the flat RS branes\\cite{RS1,RS2}, while the positive-constant curvature branes correspond to the dS branes\\cite{GaS}. Our main concern is of course the dS brane case, but we treat the flat brane case simultaneously to make clear the similarities and differences between the two cases.  The dS brane case is of particular interest because it gives a good model of braneworld inflation. In the standard 4-dimensional inflation, the quantum vacuum fluctuations play a very important role. It is therefore natural to ask if the quantum radion fluctuations play an important role, if not disastrous, in the braneworld inflation.  It should be mentioned that historically a very similar situation was analyzed by Garriga and Vilenkin \\cite{GV91,GV92} in which they considered the fluctuations of a thin domain wall in $(N+1)$ spacetime dimensions. Although they assumed the Minkowski background, many of the results obtained there apply equally to the present case. In particular, they showed that the wall fluctuation mode is represented by a scalar field living on the $N$-dimensional de Sitter space which describes the internal metric on the domain wall, and the scalar field has the negative mass squared $-NH^2$.  The paper is organized as follows. In Sec.~\\ref{sec:bg}, we describe the background spacetime and our notation. In Sec.~\\ref{sec:pert}, we solve the perturbation equation in the bulk that describes the radion mode. We find there is a gauge degree of freedom that should be carefully treated in the dS brane case in contrast to the flat brane case where no such subtlety arises. In Sec.~\\ref{sec:qradion}, assuming the dependence on the extra dimensional coordinate that solves the perturbation equation in the bulk, we derive the effective action for the radion and quantize it. In Sec.~\\ref{sec:omega}, based on the result obtained in Sec.~\\ref{sec:qradion}, we evaluate the effective energy density of the radion on the brane which is present at the linear order in the radion amplitude, and estimate its effect by calculating the rms value. We find the effect remains small for reasonable values of the model parameters even though the radion itself is unstable. In Sec.~\\ref{sec:summary}, we summarize our results and discuss the implications. ",
        "conclusions": "\\label{sec:summary}  In this paper, taking up the Randall-Sundrum type two-brane scenario in an $(n+1)$-dimensional spacetime, we have investigated the quantum fluctuations of the relative displacement of the branes, which we called the quantum radion. We have considered the cases of the flat two-brane and de Sitter two-brane systems simultaneously. Adopting the so-called Randall-Sundrum gauge, we have first solved the linear gravitational perturbation equations that describe the radion mode.  Then assuming the perturbation of the form with a fixed $z$-dependence that solves the gravitational equation in the $z$-direction, where $z$ is the extra-dimensional coordinate orthogonal to the branes, we have derived the effective action for the radion. With the effective action at hand, we have quantized the radion assuming the radion state is the Bunch-Davis vacuum.  We have analyzed the effect of the quantum radion on the brane using the effective Einstein equations derived by Shiromizu, Maeda and Sasaki\\cite{SMS} in which the radion adds an effective energy-momentum tensor at the linear order in the field amplitude. Although the radion has the negative mass squared $-nH^2$ on the de Sitter brane where $H$ is the Hubble parameter, we have found that the corresponding instability does not show up on the brane. We have noted, however, that the anisotropic stress induced by the radion is unusually large, though it still decays in time as $a^{-1}$ for a fixed $k$, where $a$ is the cosmic scale factor and $k$ is the comoving wavenumber, and the rms value of the effective energy density for a fixed $k$ decays as $a^{-3}$ irrespective of the spacetime dimensions $n$.  Focusing on the positive tension de Sitter brane with the Hubble rate $H_+$, which models the braneworld inflation, we have estimated the rms total energy density of the quantum radion by integrating over $k$ up to $aH_+$. We have introduced the density parameter $\\Omega_E$ which describes the relative magnitude of the radion energy density to the background energy density, and discussed the condition $\\Omega_E\\ll1$ on the model parameters. For $H_+\\ell_{\\rm pl}\\ll1$ and $\\ell_5\\lesssim \\ell$, which are reasonable to assume for the background to be not in the quantum regime, we have found that practically any choice of the location of the negative tension brane is allowed. This implies the quantum radion does not seriously affect the braneworld inflation scenario, at least at the linear perturbation order.  The most intriguing question remained now is if the analysis here at the order of the linear perturbation level is sufficient. Naively, one may regard the effective action of the radion we have obtained as a piece to be added to the total effective action for an effective 4-dimensional theory that includes the gravity. Then the variation of the radion effective action with respect to the 4-metric will give the energy-momentum tensor of a scalar field with a negative mass squared, which would grow and diverge in contrast to the decaying effective energy-momentum tensor we analysed in this paper. However, such an energy-momentum tensor is quadratic in the radion, which is beyond the accuracy of our analysis at the linear-order perturbation level.  We do need to investigate the higher order perturbation to truly see the effect of the quantum radion on the brane. This issue is left for future study."
    },
    "0201/astro-ph0201062_arXiv.txt": {
        "abstract": "The interaction of high-velocity neutron stars with the interstellar medium produces bow shock nebulae, where the relativistic neutron star wind is confined by ram pressure.  We present multi-wavelength observations of the Guitar Nebula, including narrow-band \\Halpha\\ imaging with HST/WFPC2, which resolves the head of the bow shock. The HST observations are used to fit for the inclination of the pulsar velocity vector to the line of sight, and to determine the combination of spindown energy loss, velocity, and ambient density that sets the scale of the bow shock.  We find that the velocity vector is most likely in the plane of the sky.  We use the \\GN\\ and other observed neutron star bow shocks to test scaling laws for their size and \\Halpha\\ emission, discuss their prevalence, and present criteria for their detectability in targeted searches.  The set of \\Halpha\\ bow shocks shows remarkable consistency, in spite of the expected variation in ambient densities and orientations.  Together, they support the assumption that a pulsar's spindown energy losses are carried away by a relativistic wind that is indistinguishable from being isotropic.  Comparison of \\Halpha\\ bow shocks with X-ray and nonthermal, radio-synchrotron bow shocks produced by neutron stars indicates that the overall shape and scaling is consistent with the same physics.  It also appears that nonthermal radio emission and \\Halpha\\ emission are mutually exclusive in the known objects and perhaps in all objects. ",
        "introduction": "\\label{Sec:intro}  Bow shocks are observed on a wide variety of astrophysical scales, ranging from planetary magnetospheres \\citep{bowPlanet} and OB-runaway stars \\citep{bowOB} to merging galaxy clusters \\citep{bowGC}.  They have been invoked in models for beamed gamma-ray bursts \\citep{bowGRB}, protostellar outflows \\citep{OLSM01}, ultra-compact \\ion{H}{2} regions \\citep{bowHII} and the interaction of the solar wind with the interstellar medium \\citep{BKK71}.  Some of the most spectacular bow shock nebulae are those associated with neutron stars. The spindown of neutron stars (NS) transfers rotational kinetic energy into the interstellar medium, and bow shock nebulae provide a way to probe these energetic environments, revealing the properties of both the NS relativistic winds and the ambient interstellar medium (ISM).  The Guitar Nebula is one such visually striking bow shock nebula, produced by the radio pulsar B2224+65. The nebula was discovered in \\Halpha\\ observations using the 5-m Hale Telescope at Palomar Observatory \\citep{CRL93}.  Until recently, there were only two other NS with bow shocks observed in \\Halpha: B1957+20 \\citep{KH88} and B0437$-$4715 \\citep{BBM+95}.  However, \\Halpha\\ bow shocks have recently been discovered around a radio quiet NS, RX~J1856.5$-$3754 \\citep{VK01}, as well as another ordinary radio pulsar, B0740$-$28 \\citep{JSG02}.  More such nebulae will no doubt be discovered as targeted searches become more sensitive and cover more objects. Additionally, there are some bow shock nebulae which are detected at radio/X-ray wavelengths, but not in \\Halpha, associated with known radio pulsars as well as radio-quiet NS: for example, the ``Duck'' \\citep{FK91} is associated with the pulsar B1757$-$24 and the supernova remnant (SNR) G5.4$-$1.2, while the cometary nebula in the SNR IC~443 is produced by a radio quiet NS \\citep{OCW+01}.   In this paper, we derive scaling relationships for bow shock parameters based on the underlying physical processes, and discuss the requirements for such nebulae to be detectable.  A detailed analysis of the Guitar Nebula is presented based on multi-wavelength observations, including high resolution Hubble Space Telescope (HST) observations which resolve the fine structure of the bow shock. Through model fitting, constraints are obtained on the density of the ambient ISM and the inclination of the NS velocity vector to the line of sight (LOS).  The published data on NS bow shocks are consolidated with the parameters for the Guitar Nebula in order to test the derived scaling relationships for \\Halpha\\ bow shocks.  We verify that the stand-off angle for plerionic bow shock nebulae detected at radio and X-ray wavelengths also scales in the same manner.  In testing these scaling laws, it is apparent that the current state of available information is inadequate, especially for radio-detected nebulae, leading to constraints on the ISM that are not very strong.  We identify the sources of uncertainty, and suggest future observations that will be able to refine current constraints as well as confirm or refute physical scenarios relating to the formation of \\Halpha\\ and plerionic bow shock nebulae.  The paper is organized as follows: the expected scaling relationships for bow shock parameters are derived in \\Sref{scale}, along with an analysis of the detectability criteria for bow shock nebulae. Multi-wavelength observations of the \\GN\\ are described in \\Sref{GN}. Since the bright head of the \\GN\\ is not resolved by ground-based optical images, high resolution HST observations are required to obtain bow shock parameters; these observations, and the associated model fitting, are described in \\Sref{HST}. In \\Sref{test}, the derived scaling laws are tested for known NS bow shocks, with consolidated parameters for both \\Halpha\\ and radio/X-ray bow shocks. Finally, in \\Sref{end}, we summarize our results, discuss their implications and identify future lines of inquiry. ",
        "conclusions": "\\label{Sec:end}  Bow shock nebulae, produced by NS moving supersonically through the ISM, constitute unique probes of the energetic environment where NS relativistic winds interact with the surrounding ISM. In this work, we report multiwavelength observations of the Guitar Nebula, a spectacular \\Halpha\\ bow shock nebula produced by a high velocity, modest \\Edot\\ radio pulsar.  High resolution HST observations show the detailed structure of the bright head of the nebula, confirming its limb-brightened nature and demonstrating the existence of a significant asymmetric component. The complex shape is best explained by the existence of a density gradient in the medium which is not aligned with the NS velocity vector.  A simple momentum conserving bow shock model is used to fit the HST observations and extract the stand-off angle $\\theta_0$, as well as placing joint constraints on the inclination to the LOS, the density of the ambient medium, and the velocity and distance of the pulsar. While the low S/N ratio of the current HST image precludes a more detailed analysis, future HST data will provide more definitive constraints for some of these parameters.  The stand-off angle $\\theta_0$ and the \\Halpha\\ flux $F_\\alpha$ scale with other measurable parameters (\\Edot, $D$, $\\mu$) and unknowns ($\\nH$, $X$). These scaling relationships are derived from the basic physics of the bow shock formation.  Using the observations of the Guitar Nebula (described here) along with published parameters for the other known \\Halpha\\ bow shocks, we confirm the scaling of $\\theta_0$ with the other parameters for a diverse collection of nebulae, from simple canonical bow shocks to contorted shapes.  The scaling relationship is used to derive upper limits on the density of the ambient medium, independent of the inclination of the bow shock to the LOS (which is hard to measure in a model-independent way). Additionally, the scaling of \\Halpha\\ flux is demonstrated with different combinations of measurable parameters: while the relationship is not as constraining in this case, the observations conform to the expected scaling within the inherent uncertainties in the neutral fraction and extinction and the unknown density of the ambient medium.  We have also consolidated the available parameters for radio and X-ray bow shocks described in the literature, both for standard radio pulsars and radio-quiet or undetected NS. The stand-off angle appears to scale as expected for these objects as well.  However, as described by \\citet{GF00} and illustrated by \\Fref{scale-radio}, the information available for these NS is inadequate, and has large margins of uncertainty.  VLBI measurements of the parallaxes and proper motions of these objects will be essential for a clearer understanding of the physics of plerionic bow shock nebulae and their correlation with young NS and well-defined supernova remnants.  The observed anti-correlation of \\Halpha\\ and plerionic radio emission from bow shocks is suggestive, though it needs to be confirmed in a larger sample of objects.  Some young NS may be inside superbubbles created by previous supernovae, where the neutral fraction is insignificant.  Pre-ionization of the medium, either by the young NS, or by the supernova in which it was born, may also suppress \\Halpha\\ emission by reducing the enutral fraction, leading to the prediction that young NS with associated supernova remnants should not produce \\Halpha\\ bowshocks.  Conversely, if radiation from a young NS pre-ionizes the medium, \\Halpha\\ bow shocks should only be associated with older NS.  This would be consistent with the apparent drop off in the efficiency with which \\Edot\\ is converted to radio emission in pulsar wind nebulae as a function of age \\citep{GSF+00}.  Thus there are multiple avenues of future enquiry that promise to be fruitful.  As the recently reported discovery of a bow shock nebula for B0740$-$28 demonstrates \\citep{JSG02}, deep observations of suitably selected pulsars will probably result in the detection of more \\Halpha\\ bow shocks, leading to constraints on the NS and ISM properties, as described in this work.  Future radio and X-ray observations will also produce more such objects, and clarify the relationship between the NS age and environment and the type of bow shock emission.  Additionally, further interferometric proper motions (and parallaxes, where possible) are required to produce firmer constraints on the properties of the ISM and the NS using known bow shock nebulae, while theoretical models of bow shock evolution can elucidate the differences in the behavior of bow shocks downstream from the stand-off region.  It is also possible that old high-velocity neutron stars, including some born in the Galactic halo, may be identified through the serendipitous discovery of radio, X-ray or \\Halpha\\ bow shock nebulae.  The Guitar Nebula, specifically, will continue to be of interest, both from a theoretical modeling perspective, and for observations of its evolution over time.  Monitoring of the pulsar dispersion measure and changes in the shape of the head of the Guitar in multi-epoch high resolution observations will provide a unique perspective on the small-scale inhomogeneities in the interstellar medium."
    },
    "0201/astro-ph0201548_arXiv.txt": {
        "abstract": "Using the method of separation of variables and a new approach to calculations of the prolate spheroidal wave functions, we study the optical properties of very elongated (cigar-like)  spheroidal particles.  A comparison of extinction efficiency factors of prolate spheroids and infinitely long circular cylinders is made. For the normal and oblique incidence of radiation, the efficiency factors for spheroids converge to some limiting values with an increasing aspect ratio $a/b$ provided particles of the same thickness are considered. These values are close to, but do not coincide with the factors for infinite cylinders. The relative difference between factors for infinite cylinders and elongated spheroids ($a/b \\ga 5$) usually does not exceed 20\\,\\% if the following approximate relation between the angle of incidence $\\alpha~({\\rm in~degrees})$ and the particle refractive index $m=n+ki$ takes the place: $\\alpha  \\ga 50 |m-1| + 5$ where $1.2 \\la n \\la 2.0$ and $k \\la 0.1$.  We show that the quasistatic approximation can be well used for very elongated optically soft spheroids of large sizes. ",
        "introduction": "Rapid calculations of light scattering by non-spherical particles are very important in many scientific and engineering applications (see discussion in \\cite{bh83}, \\cite{mht00}). The simplest model of non-spherical particles --- an infinitely long circular cylinder is not physically reasonable. However, it looks attractive to find cases when this model could be useful because the calculations for infinite cylinders are very simple. Therefore, we compare the light scattering by {\\it elongated spheroids} and {\\it infinite cylinders}. Previously, such a comparison was made by Martin \\cite{mar78} and Voshchinnikov \\cite{v90b} for normal incidence of radiation (perpendicular to the rotation axis of a particle, $\\alpha=90^)$).  Our consideration is based on the solution to the light scattering problem for spheroidal particles by the {\\it Separation of Variables Method} (SVM) (see \\cite{f83}, \\cite{vf93} for details). A new type of expansions of the prolate wave functions (Jaff\\'e expansions  \\cite{fv01}) opens a possibility to calculate the optical properties of very elongated (cigar-like) particles.  In this paper, we study the optical properties of prolate homogeneous spheroids with large aspect ratios and compare them with those of infinite circular cylinders and of spheroids in quasistatic approximation. ",
        "conclusions": "We applied a new method of calculations of the spheroidal wave functions to the study of the optical properties of very elongated (cigar-like)  spheroids. New approach allowed to compare light scattering by prolate spheroids and infinitely long circular cylinders and investigate the applicability of  the quasistatic approximation for very elongated spheroids.  It is found that the  efficiency factors for spheroids and cylinders have quite similar behaviour for the normal and oblique incidence of radiation, if the aspect ratio of spheroids $a/b \\ga 5$. The resemblance of factors arises provided spheroids and very long cylinders of the same volume and aspect ratio are considered. The following approximate relation between the angle of incidence $\\alpha~({\\rm in~degrees})$ and the particle refractive index $m=n+ki$ takes the place: $\\alpha  \\ga 50 |m-1| + 5$ where $1.2 \\la n \\la 2.0$ and $k \\la 0.1$. In this case, the relative discrepancy between factors for infinite cylinders and elongated spheroids ($a/b \\ga 5$)  does not exceed 20\\,\\% near the  first maximum.  It is shown that the quasistatic approximation rather well describes the extinction by very elongated optically soft spheroids of large sizes.   \\ack{The authors are thankful to Vladimir Il'in for useful comments. The work was partly supported by the INTAS grant 99/652, grant for scientific school on theoretical astrophysics, and the program ``Astronomy'' of the Russian Federal Government.}"
    },
    "0201/astro-ph0201254_arXiv.txt": {
        "abstract": "{ We have compiled optical and radio astrometric data of the microquasar \\object{LS~5039} and derived its proper motion. This, together with the distance and radial velocity of the system, allows us to state that this source is escaping from its own regional standard of rest, with a total systemic velocity of about $150$~km~s$^{-1}$ and a component perpendicular to the galactic plane larger than 100~km~s$^{-1}$. This is probably the result of an acceleration obtained during the supernova event that created the compact object in this binary system. We have computed the trajectory of \\object{LS~5039} in the past, and searched for OB associations and supernova remnants in its path. In particular, we have studied the possible association between \\object{LS~5039} and the supernova remnant \\object{G016.8$-$01.1}, which, despite our efforts, remains dubious. We have also discovered and studied an \\ion{H}{i} cavity in the ISM, which could have been created by the stellar wind of \\object{LS~5039} or by the progenitor of the compact object in the system. Finally, in the symmetric supernova explosion scenario, we estimate that at least $17\\,M_{\\sun}$ were lost in order to produce the high eccentricity observed. Such a mass loss could also explain the observed runaway velocity of the microquasar. ",
        "introduction": "\\label{introduction}  Microquasars are stellar-mass black holes or neutron stars that mimic, on smaller scales, many of the phenomena seen in AGN and quasars. These objects have been found in X-ray binary systems, where a compact object accretes matter from a companion star. Radio emitting X-ray binaries with relativistic radio jets, like \\object{SS~433}, \\object{GRS~1915+105}, \\object{GRO~J1655$-$40} or \\object{Cygnus~X-3}, are good examples of microquasars (see Mirabel \\& Rodr\\'{\\i}guez \\cite{mirabel99} for a detailed review). With the recent addition of \\object{LS~5039} (Paredes et~al. \\cite{paredes00}), \\object{Cygnus~X-1} (Stirling et~al. \\cite{stirling01}) and \\object{XTE~J1550$-$564} (Hannikainen et~al. \\cite{hannikainen01}) to the microquasar group, the current number of this kind of sources is 14, among $\\sim280$ known X-ray binaries (Liu et~al. \\cite{liu00}; Liu et~al. \\cite{liu01}) of which $\\sim50$ display radio emission. Recent studies of microquasars can be found in Castro-Tirado et~al. (\\cite{castro01}).  Attention to the star \\object{LS~5039} was first called by Motch et~al. (\\cite{motch97}), who proposed it as a High Mass X-ray Binary (HMXB) candidate associated with the X-ray source \\object{RX~J1826.2$-$1450}. The object is located at an estimated distance of $\\sim3.1$~kpc and close to the galactic plane ($l=16.88\\degr$, $b=-1.29\\degr$). Soon after, non-thermal and moderately variable radio emission was reported by Mart\\'{\\i} et~al. (\\cite{marti98}) using the Very Large Array (VLA). The evidence of its microquasar nature was provided by Paredes et~al. (\\cite{paredes00}) when radio jets were discovered with the Very Long Baseline Array (VLBA) at milliarcsecond (mas) angular scales. These authors also pointed out the possible connection of \\object{LS~5039} with \\object{3EG~J1824$-$1514}, i.e., one of the unidentified EGRET sources of high energy $\\gamma$-rays. X-ray observations of \\object{RX~J1826.2$-$1450} by Rib\\'o et~al. (\\cite{ribo99}) did not reveal pulsations and were consistent with a significantly hard X-ray spectrum up to 30~keV, with a strong Gaussian iron line at 6.6~keV.  The mass donor in \\object{LS~5039} was originally classified as an O7V((f)) star by Motch et~al. (\\cite{motch97}). This classification has been recently improved thanks to optical and near infrared spectroscopic observations by Clark et~al. (\\cite{clark01}), which indicate an O6.5V((f)) star.   \\begin{table*} \\begin{center} \\caption[]{Compilation of optical and radio positions, with associated errors, of \\object{LS~5039}.} \\label{positions} \\begin{tabular}{lllclcll} \\hline \\noalign{\\smallskip} Wavelength & Epoch & $\\alpha$~(ICRS) & $\\sigma_{\\alpha\\cos\\delta}$ & $\\delta$~(ICRS)               & $\\sigma_\\delta$ & Catalog & Reference\\\\ Domain &      & (h, m, s)       & (mas)                       & ($\\degr$, $\\arcmin$, $\\arcsec$) & (mas)           & name    &\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} Optical & 1905.45  & 18 26 15.0194 & 255 & -14 50 53.300 & 247 & AC~2000.2 & Urban, private communication\\\\ & 1907.43  & 18 26 15.0177 & 255 & -14 50 53.075 & 247 & AC~2000.2 & Urban, private communication\\\\ & 1951.577 & 18 26 15.034  & 250 & -14 50 53.59  & 250 & USNO-A2.0 & Monet et~al. \\cite{monet99}\\\\ & 1979.484 & 18 26 15.0557 & ~64 & -14 50 54.075 & ~47 & TAC~2.0   & Zacharias \\& Zacharias \\cite{zacharias99a}\\\\ & 1986.653 & 18 26 15.054  & 300 & -14 50 54.29  & 300 & GSC~1.2   & Morrison et~al. \\cite{morrison01}\\\\ & 1991.75  & 18 26 15.0427 & 149 & -14 50 54.229 & 120 & Tycho-2   & H{\\o}g et~al. \\cite{hog00}\\\\ & 2000.289 & 18 26 15.0563 & ~13 & -14 50 54.277 & ~13 & UCAC1     & Zacharias et~al. \\cite{zacharias00}\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} Radio   & 1998.24  & 18 26 15.056  & ~10 & -14 50 54.24  & ~10 & VLA obs.  & Mart\\'{\\i} et~al. \\cite{marti98}\\\\ & 2000.42  & 18 26 15.0566 & ~~4 & -14 50 54.261 & ~~6 & VLBA obs. & Rib\\'o et~al. \\cite{ribo02}\\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\end{center} \\end{table*}   Our present knowledge of the system orbit is based only on the radial velocity measurements in the optical by McSwain et~al. (\\cite{mcswain01}). Their most remarkable findings consist of a short orbital period of $P=4.117$~days and a significant eccentricity of $e=0.41$. They were also able to determine the system radial velocity and a mass function of $f(m)=0.00103\\,M_{\\sun}$. The fact that the available optical photometry does not show an ellipsoidal modulation could indicate that the orbital inclination is very low. Therefore, the possibility of having a black hole in the system cannot be ruled out in spite of the small mass function observed.  The formation of the compact object in an X-ray binary necessarily requires a supernova explosion that does not disrupt the system. This explosive event is expected to considerably change the kinematic properties of the binary star. The change may be rather extreme (i.e., kick velocities approaching $\\sim10^3$ km~s$^{-1}$) for highly asymmetric supernovae, as it has been proposed for the microquasar \\object{Circinus~X-1} (Tauris et~al. \\cite{tauris99}). This mechanism may also be responsible for ejecting X-ray binaries into the halo of the Galaxy. The fast moving X-ray nova \\object{XTE~J1118+480} is a possible example (Mirabel et~al. \\cite{mirabel01}), although it could have been ejected from a globular cluster in the past. On the other hand, the microquasars \\object{GRO~J1655$-$40} and \\object{Cygnus~X-1}, also display runaway velocities (Shahbaz et~al. \\cite{shahbaz99}; Kaper et~al. \\cite{kaper99}). With independence of their origin in the galactic plane or in the halo, the existence of runaway microquasars is a new and important issue that deserves an in-depth study.  In this paper we focus on the kinematic properties of \\object{LS~5039} and its surroundings. In Sect.~\\ref{ppm} we estimate the proper motions of the system, using optical and radio data. In Sect.~\\ref{ejection} we discuss the distance and space velocity of \\object{LS~5039}, whereas in Sect.~\\ref{trajectory} we analyze the trajectory of the binary system in the past. In Sect.~\\ref{snr} we study the possible association of this microquasar with a supernova remnant, and in Sect.~\\ref{surroundings} we analyze their  \\ion{H}{i} surroundings. Finally, we make a global discussion in Sect.~\\ref{discussion} and summarize our main conclusions in Sect.~\\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  After an in-depth study of the proper motions and surroundings of \\object{LS~5039} our main conclusions are:  \\begin{enumerate}  \\item Positions at optical and radio wavelengths have been used to compute independent optical and radio proper motions, which are perfectly compatible. Therefore, based only on astrometric data we are able to confirm that both, the optical and the radio emission, originate in the same object.  \\item From the combined optical and radio positions we have computed an accurate proper motion for \\object{LS~5039}. This, together with the new estimate of $2.9\\pm0.3$~kpc for the distance, allows us to compute a space velocity of ($U=51$, $V=-71$, $W=-118$) km~s$^{-1}$ in its Regional Standard of Rest (RSR). This results implies that \\object{LS~5039} is a runaway microquasar with $v_{\\rm sys}\\simeq150$ km~s$^{-1}$, escaping from its own RSR with a large velocity component perpendicular to the galactic plane. This is probably the result of the SN event that created the compact object in this binary system.  \\item We have computed the past trajectory of \\object{LS~5039}. Two OB associations have been found close to its path in the plane of the sky. However, they are too close to us to be related to the microquasar. On the other hand, we have also found three SNRs near the path of \\object{LS~5039}. After discarding two of them based on distance arguments, we have focused our attention on \\object{SNR G016.8$-$01.1}. A study of this source could not clearly confirm nor reject the association due to the large uncertainties in the estimated radio flux density of the SNR. This fact perhaps justifies future, high sensitivity searches of low-brightness remnants in this region.  \\item We have found a semi-open \\ion{H}{i} cavity close to the \\object{LS~5039} position. Although the O((f)) star in this microquasar seems to be the main agent forming the bubble, a contribution from the progenitor of the compact object cannot be ruled out.  \\item Finally, we are able to explain both, the high space velocity and the high eccentricity observed, in a symmetric SN explosion scenario with a mass loss of $\\Delta\\,M\\sim17\\,M_{\\sun}$.  \\end{enumerate}"
    },
    "0201/astro-ph0201124_arXiv.txt": {
        "abstract": "We use N-body simulations to test the predictions of the redshift distortion in the power spectrum given by the halo model in which the clustering of dark matter particles is considered as a result both of the clustering of dark halos in space and of the distribution of dark matter particles in individual dark halo. The predicted redshift distortion depends sensitively on several model parameters in a way different from the real-space power spectrum. An accurate model of the redshift distortion can be constructed if the following properties of the halo population are modelled accurately: the mass function of dark halos, the velocity dispersion among dark halos, and the non-linear nature of halo bias on small scales. The model can be readily applied to interpreting the clustering properties and velocity dispersion of different populations of galaxies once a cluster-weighted bias (or equivalently an halo occupation number model) is specified for the galaxies. Some non-trivial bias features observed from redshift surveys of optical galaxies and of IRAS galaxies relative to the standard low-density cold dark matter model can be easily explained in the cluster weighted bias model. The halo model further indicates that a linear bias can be a good approximation only on for $k \\leq 0.1 hMpc^{-1}$. ",
        "introduction": "The power spectrum of the galaxy spatial distribution is an important statistic for describing inhomogeneities in the Universe. The spatial distribution of galaxies observed with a redshift survey is distorted by the peculiar motions of galaxies, and a statistically isotropic distribution (e.g. the power spectrum or the correlation function) in real space becomes anisotropic in redshift space (Geller \\& Peebles 1973; Davis \\& Peebles 1983; Bean et al. 1983; Kaiser 1987). On the other hand, when measuring clustering for high redshift objects in redshift space, choosing the wrong cosmological model can lead to an additional anisotropy in the redshift space distribution (Matsubara \\& Suto 1996; Ballinger, Peacock \\& Heavens 1996). The theory for the redshift distortions caused by an assumed world model and by the large scale linear motions is now well established. The transform of the clustering pattern from the true world model to an assumed model is just a simple mapping in the coordinates. On large scales and for a linear bias, the redshift power spectrum $P^{S}_{l}(k, \\mu)$ can be derived (Kaiser 1987) \\begin{equation} P^{S}_{l}(k, \\mu) = P^{R}_{l}(k)\\, [1+\\beta \\mu^{2}]^{2} \\label{Pk} \\end{equation} where $\\mu$ is the cosine of the angle between the line-of-sight and the $\\bf k$ \\rm vector, $\\beta=\\Omega^{0.6}/b$, $\\Omega$ the density parameter, $b$ the linear bias, and $P^{R}_{l}(k)$ the linear power spectrum of dark matter in real space. The hope has been to measure the dynamical quantity $\\beta$ and the cosmological parameters through studying the anisotropy of galaxies or galaxy cluster on large enough scales. However, it has been shown that the virialized motions within rich clusters (the Finger-of-God effect) are so prominent that the clustering pattern on large scales of wave number $k\\sim 0.1\\impc$, is significantly affected (Cole, Fisher \\& Weinberg 1994; Suto et al. 1999). The effect of non-linear motions on the redshift distortion must be properly modelled in order to measure the dynamical quantity and the cosmological parameters from the redshift distortion of extragalactic objects. Recent results (Peacock et al. 2001) from the redshift distortion of the 2dF galaxies further support this point, as even with the largest redshift survey available today there exists a tight degeneracy in the determinations of the parameters $\\beta$ and the pairwise velocity $\\sigma_v$ (reflecting the non-linear motion).  As pointed out recently by Jing \\& B\\\"orner (2001, hereafter JB2001), the existing analytical model for the nonlinear velocity distortion, e.g. the exponential distribution of the relative velocity in coordinate space (Davis \\& Peebles 1983) or the Lorentz damping function in k-space (Cole, Fisher \\& Weinberg 1995; Peacock \\& Dodds 1994), is at best an approximation for some scales. Although JB2001 have made an extensive study of the redshift distortion for the dark matter and for two specific biased tracers in three typical cosmological models based on high-resolution simulations, it is unknown how to generalize their results to cosmological models and/or to biased tracers different from those studied in their work without running new simulations. In this paper, we present an analytical model for the redshift distortion based on the halo model, and we will test the accuracy of the analytic model with the results of JB2001.  A key concept in the standard hierarchical scenario of structure formation is the formation of dark matter halos, which are virialized systems of dark matter particles formed through non-linear gravitational collapse in the cosmic density field. Since the formation of dark halos involves only gravitational physics, accurate analytic models are now available for many properties of the halo population, such as the mass function (Press \\& Schechter 1974; Lee \\& Shandarin 1998; Sheth \\& Tormen 1999; Sheth, Mo \\& Tormen 2001), clustering properties, (Mo \\& White 1996; Mo, Jing \\& White 1997; Jing 1998, 1999; Sheth \\& Tormen 1999; Sheth, Mo \\& Tormen 2001; Hamana et al. 2001), and density profiles (Navarro, Frenk \\& White 1996, 1997; Moore et al. 1999; Jing \\& Suto 2000; Klypin et al. 2001).  Such models are very useful in understanding the clustering properties of matter in the universe, as well as in understanding the bias of the galaxy distribution relative to the underlying density field. Indeed, since in the hierarchical cosmogony all masses in the universe are partitioned in dark halos, the halo clustering properties , the halo mass function and the halo density profiles are  sufficient for the construction of a clustering model for the dark matter in the universe (Scherrer \\& Bertschinger 1991; Sheth \\& Jain 1997; Ma \\& Fry, 2000a,b; Seljak 2000; Cooray, Hu \\&  Miralda-Escude 2000; Cooray 2001). Furthermore, since galaxies are assumed to form through gas cooling and condensation in dark halos, a halo  model of galaxy clustering can also be constructed by combining the  clustering properties of the halo population with an assumption of  how galaxies populate dark halos (Jing, Mo \\& B\\\"orner 1998; Peacock \\& Smith 2000; Seljak 2000; Scocimarro \\& Sheth 2001; Scoccimarro et al. 2001; Berlind \\& Weinberg 2001).  The halo model also provides a useful way to understand the redshift distortion in galaxy clustering (Seljak 2001; White 2001). In this model, the enhancement of the redshift-space power spectrum (relative to that in real space) on large scales is assumed to arise in the halo-halo correlation, while the smearing (Finger-of-God) effect on small scales is attributed to the velocity dispersions in dark halos. In this paper we show, however, that the redshift distortion depends on several important effects which were not included in early modelling.  Using results derived from high-resolution N-body simulations we show  that an accurate model for the redshift distortion dependents not only on the mass function of dark halos, the dark halo clustering and the velocity dispersion among dark halos, but also on the nonlinear motion of halos at intermediate scales. ",
        "conclusions": "\\begin{figure} \\centering \\vskip-0.7cm \\mbox{\\psfig{figure=fig10.ps,height=8.8cm,width=8.2cm}} \\caption{The same as Figure 9, but for SCDM model.} \\end{figure}  In this paper we present an analytical model for the non-linear redshift-space power spectrum of dark matter and of galaxies based on a halo prescription. The model has three important ingredients: the halo mass function, the mass density profile of halos, and the halo-halo redshift space power spectrum. The predicted redshift power spectrum is found to be insensitive to the details of the halo density profiles: the NFW density profile and a steeper inner density profile (Jing \\& Suto 2000) yield an indistinguishable redshift power spectrum. When we use, as many others have done (for the real space correlation function of dark matter and for the radial averaged redshift power spectrum) the Press-Schechter formula for the halo mass function, the Mo \\& White formula for the halo bias, and Kaiser's formula for the redshift distortion of the halo spatial distribution, we find that the predicted redshift-space power spectrum for dark matter is too high at large $\\mu$ to be consistent with the high-resolution simulation results of JB2001. The reason why the halo model works well for the real space clustering but not for the redshift-space power spectrum is that in the latter case the result is dominated by the two-halo term on non-linear scales.  We have checked carefully which assumptions have caused the discrepancy between the halo model prediction and the simulations. First the fitting formulae of the halo mass function and the halo bias from numerical simulations (Sheth \\& Torman 1999) are used in the replacement of the analytical formulae based on the Press-Schechter formalism.  This deteriorates the agreement between the halo model prediction and the simulation results , because there are more small halos and the spatial clustering of small halos is stronger in the modified formulae. We also found that the bias of the halo distribution is nearly linear relative to the linear density field, so the non-linearity of the bias could not be the main contribution to the discrepancy. Instead we found that the non-linear motions of the halos, which were neglected in previous studies, are the main cause. Once we take this effect into account the redshift power spectrum based on the halo model agrees very well with the simulation results. Furthermore, the redshift space power spectrum can be precisely predicted if the halo occupation number model, e.g. the cluster weighted model, is given for the galaxies.  Our results show for the first time that the two-halo term can dominate some statistics in the redshift space even at small scales. Therefore a halo model based on the density profiles of halos and on the redshift distribution of halos predicted by the {\\it linear theory} may become inaccurate for some statistics of redshift clustering. For the time being, an accurate model (analytical or fitting) for the redshift power spectrum of the halos, which includes the effects of the nonlinear bias and the non-linear motions, is needed for the prediction of the redshift space power spectra of galaxies and dark matter.  In combination with the cluster weighted bias model, we show that a non-trivial scale-dependent bias is generally expected for galaxies in CDM models. The bias could be regarded linear only on the scales at the wavelength larger than $60\\mpc$ (also see Seljak 2001). The bias is different for different populations of galaxies as observed. The currently favored LCDM model can well explain the observed features of the spatial bias reported for optical galaxies and for IRAS galaxies recently, if the cluster weighted bias model is applied."
    },
    "0201/astro-ph0201312_arXiv.txt": {
        "abstract": "The standard calculations of the synchrotron emission from charged particles in magnetic fields does not apply when the energy losses of the particles are so severe that their energy is appreciably degraded during one Larmor rotation. In these conditions, the intensity and spectrum of the emitted radiation depend on the observation time $T_{obs}$: the standard result is recovered only in the limit $T_{obs}\\ll T_{loss}$, where $T_{loss}$ is the time for synchrotron losses. In this case the effects of the radiation backreaction cannot be detected by the observer. We calculate the emitted power of the radiation in the most general case, naturally including both the cases in which the backreaction is relevant and the standard case, where the usual result is recovered. Finally we propose several scenarios of astrophysical interest in which the effects of the backreaction cannot and should not be ignored. ",
        "introduction": "As pointed out in \\cite{syn1} (hereafter paper I), some aspects of the synchrotron emission of high energy particles did not receive proper attention in the literature, mainly because the standard treatment proved to be valid over most of the energy range of interest for astrophysical applications. Recently, this range of interest changed considerably, mainly due to the discovery of radiation processes at ultra-high energies. Nevertheless the important corrections to synchrotron emission in this regime have been ignored and the standard calculations have been adopted.  In paper I we explored the generalization of the calculations of the synchrotron emission from an ensemble of particles radiating coherently. We considered there the cases of bunches of particles both in the monoenergetic case and in the case of a spectrum of particles, and for each we established the criteria for the synchrotron emission to be coherent.  In the present paper we explore an effect that becomes relevant at sufficiently high energies, when the energy lost by a particle during one Larmor gyration becomes comparable with the energy of the particle itself. We call this effect backreaction, with may be a slightly improper term. We find that in this regime there are important corrections to the spectra of the emitted radiation in numerous scenarios currently discussed in the literature, confirming but extending previous findings of Ref. \\cite{nelson}, where the case of a monoenergetic distribution of radiating electrons was considered. In this paper we also study the situation of an ensemble of particles with arbitrary energy spectra and find an approximate analytical solution of the problem of the synchrotron backreaction, that may be useful to estimate the magnitude of the effect. The formalism used here, introduced in paper I and briefly summarized in this paper, allows us to take naturally into account possible coherence effects in the synchrotron radiation, together with the backreaction. Cases of astrophysical interest in which the effects of the backreaction are supposed to play an important role will be discussed.  The paper is structured as follows: in section 2 we describe the backreaction and the effects that can be expected on simple basis. In section 3 we calculate the spectrum of the radiation emitted by monoenergetic particles. In section 4 we generalize the calculation to the case of a spectrum of radiating particles and we present analytical approximations that allow us to estimate the effects of the backreaction without (or before) being involved in the detailed calculations. We present our conclusions in section 5. ",
        "conclusions": "We calculated the spectrum of the synchrotron emission from a system of charged particles in the most general case in which the observation time is arbitrary compared to the time for the energy losses of the particles. We first calculated the effect for the case of a particle with Lorentz factor $\\gamma_0$.  Our findings on the monoenergetic case can be summarized as follows:  {\\it i)} There is a critical frequency $\\omega_{cr}$ such that the spectrum of the radiation at $\\omega\\le \\omega_{cr}$ is the usual synchrotron spectrum $\\propto \\omega^{1/3}$, while for $\\omega\\ge \\omega_{cr}$ the backreaction affects the spectrum changing it to $\\propto \\omega^{-1/2}$.  {\\it ii)} The critical frequency $\\omega_{cr}$ depends on the observation time but it turns out to be independent of the Lorentz factor of the radiating particles $\\gamma_0$. The dependence on these parameters is as found in eq. (\\ref{eq:ommax2}), so that increasing the observation time, power is moved to gradually lower frequencies.  {\\it iii)} The standard limit is recovered when the observation time is much smaller than the time for losses $\\omega_R T_{obs}\\ll 1$.  The more realistic case investigated in this paper is that of a power law spectrum of particles $N(\\gamma_0)\\propto \\gamma_0^{-p}$. In general what happens is that there may be a critical Lorentz factor $\\gamma_c$, such that particles with $\\gamma_0\\ge\\gamma_c$ are affected by the backreaction while the particles with $\\gamma_0\\le \\gamma_c$ behave in the standard way. The spectrum of the radiation in this case is a superposition of different components. Our findings can be summarized as follows:  {\\it a)} the spectrum of the radiation can be divided into a low frequency one and a high frequency one, with the separation occurring at the frequency $\\omega_{cr}=3\\omega_L/(\\omega_K T_{obs})^2$.  {\\it b)} the particles with Lorentz factor $\\gamma_0\\le\\gamma_c$ radiate the standard synchrotron radiation whose spectrum is $\\propto \\omega^{-(p-1)/2}$.  {\\it c)} the particles with $\\gamma_0\\ge\\gamma_c$ do radiate in regime of backreaction and affect both the low frequency and the high frequency regime.  {\\it d)} the low frequency spectrum radiated by particles with $\\gamma_0\\ge\\gamma_c$ is $\\propto \\omega^{1/3}$ for $\\omega\\ll \\omega_{cr}$, but it is always dominated, in the same frequency range, by the standard synchrotron radiation.  {\\it e)} the high frequency radiation radiated by particles with $\\gamma_0\\ge\\gamma_c$ is $\\propto \\omega^{-p/2}$. Therefore it represents a suppression of the radiation compared to the case of standard synchrotron radiation. Note that increasing the observation time, while the slopes at low and high frequency remain unchanged, the boundary between the two regimes moves toward lower frequencies, so that, as a consequence, the height of the spectrum at high frequencies becomes increasingly lower.  {\\it f)} At fixed observation time (which is obviously decided by the observer) the critical frequency only depends on the magnetic field in the production region. This is a very important point: in the standard synchrotron emission, it is in general not possible to extract the magnetic field from a measurement of the synchrotron flux because there is degeneracy between the number of radiating particles and the strength of the magnetic field. In the backreaction case however, the position of the change in slope uniquely defines the magnetic field, so that the measurement can in principle be used to directly infer the strength of the magnetic field.  The crucial question is whether there are situations in which the conditions for the backreaction to be relevant are fulfilled. The answer can be found in fig. 1, where we plotted the regions of interest for both electrons and protons in the plane $B-E$. At each magnetic field there corresponds a range of energies for which the backreaction is important. We immediately see that for conditions typical in the Galaxy, $B\\sim 1\\mu G$ only electrons with $E\\simgt 10^7$ GeV feel the effects of the backreaction. On the other hand in magnetic fields which are typical of neutron stars, $B\\sim 10^{10}-10^{12}$ Gauss, electrons with energies in excess of MeV-GeV already need to be accounted for in the frame of a backreaction approach. In the same environment, protons with energies larger than $10^5-10^6$ GeV also radiate in a regime in which backreaction is important.  Some applications of the calculations and results illustrated in this paper will be presented in forthcoming papers. They include: the synchrotron emission from ultra-high energy electron-positron pairs generated as a result of the decay of super-heavy relics \\cite{bere} in the Galaxy {\\cite{blasi} or in the decay of the $Z^0$ resonance produced in $\\nu\\bar\\nu$ annihilation (Z-burst model \\cite{weiler,fargion}). In this case we have magnetic fields $B\\sim 1\\mu G$ and energies in excess of $10^{10}$ GeV, so that the synchrotron radiation is strongly affected by the backreaction.  A more {\\it conventional} application concerns the TeV gamma ray emission from BL Lac objects. In \\cite{aharonian,protheroe} a proton synchrotron model was proposed in which the TeV emission is the result of the synchrotron emission of ultra high energy protons ($E\\sim 10^{19}$ eV) in a magnetic field of order $\\sim 100$ Gauss. We estimate that with these parameters and for observation times of the order of the duration of the observed flares (a few hours) the backreaction affects visibly the TeV gamma ray spectra \\cite{AloBla}.  {\\bf Aknowledgments} We are very grateful to F. Pacini and A. Olinto for many useful discussions and to M. Salvati and T. Stanev for a critical reading of the manuscript. We are also grateful to the anonymous referee for the interesting remarks on the paper."
    },
    "0201/astro-ph0201474_arXiv.txt": {
        "abstract": "The long white-dwarf spin periods in the magnetic cataclysmic variables EX~Hya and \\vcen\\ imply that if the systems possess accretion discs then they cannot be in equilibrium. It has been suggested that instead they are discless accretors in which the spin-up torques resulting from accretion are balanced by the ejection of part of the accretion flow back towards the secondary.   We present phase-resolved spectroscopy of \\vcen\\ aimed at deducing the nature of the accretion flow, and compare this with simulations of a discless accretor.  We find that both the conventional disc-fed model and the discless-accretor model have strengths and weaknesses, and that further work is needed before we can decide which applies to \\vcen. ",
        "introduction": "The magnetic cataclysmic variables are close binary stars in which one can study the interaction of an accretion flow with a magnetic field. Where the accreting white dwarf is only weakly magnetic (\\sqiglt 10$^{5}$ G) an accretion disc disc in a manner little different from that in non-magnetic systems. Stronger fields (\\sqiggt 10$^{7}$ G) lock the spin of the white dwarf to the binary orbit and dominate the accretion flow, forcing it to accrete along field lines. The intermediate case is less clear, and systems in this category (refered to as intermediate polars or IPs) display a range of behaviours depending on the mass-transfer rate, field strength and white-dwarf spin period.  Among the possibilities are (1) a partial disc which is disrupted when the magnetic pressure exceeds the ram pressure, giving way to magnetically channelled flow inside the magnetosphere; (2) a partial disc, but with some of the accretion stream overflowing the disc to interact directly with the magnetosphere (e.g.\\ Hellier 1991); (3) discless accretion in which the flow can be regarded as diamagnetic, intermediate between the purely ballistic and magnetically channeled flows (e.g.\\ King 1993; Wynn \\& King 1995); (4) models in which the propeller effect of a rapidly spinning field prevents accretion (e.g.\\ Wynn, King \\&\\ Horne 1997).  Recent reviews of these possibilities are presented in Hellier (2001), chapter 9, and Wynn (2001).  Of particular relevance to this paper is the ratio of the spin period of the white dwarf to the orbital period of the binary.  Most IPs have \\Pspin/\\Porb\\,\\sqiglt 0.1, and indeed no system can both possess an accretion disc and be in equilibrium unless this inequality holds.  This condition is equivalent to the condition \\rco\\,\\sqiglt \\rcirc\\ where \\rco\\ is the corotation radius (the radius at which the magnetic field corotates with a Keplerian flow) and \\rcirc\\ is the circularization radius (the radius at which a circular orbit has the same angular momentum as the stream from the inner Lagrangian point). However, King \\& Wynn (1999) discovered that a discless system can reside on a continuum of equilibria with \\rcirc\\,\\sqiglt \\rco\\,\\sqiglt $b$, where $b$ is the distance to the Lagrangian point. Such a system would have a longer spin period, with 0.1\\,\\sqiglt \\Pspin/\\Porb\\,\\sqiglt 0.7.  At the time only one IP (EX~Hya, with \\Pspin\\ = 67 mins and \\Porb\\ = 98 mins) was known to have a secure \\Pspin/\\Porb\\ ratio greater than \\appro 0.1.  The purpose of this paper is to (1) confirm earlier indications that \\vcen\\ is a second system in this category, and (2) analyse spectroscopic observations to investigate whether the accretion flow is better described by the partial-disc model or by the diamagnetic-flow model.  Buckley \\etal\\ (1998)'s discovery paper on \\vcen\\ (RX\\,J1238--38) and follow-up X-ray observations (Hellier, Beardmore \\&\\ Buckley 1998) found a spin period of 2147 s (revealed by an optical and X-ray pulsation), and suggested an orbital period near 85--90 mins, and thus a \\Pspin/\\Porb\\ ratio of \\appro 0.4.  The star is also notable for showing a 1860-s optical and $J$-band periodicity (Buckley \\etal\\ 1998). Given the above spin and orbital periods, the only plausible identification is with the first harmonic of the beat cycle between the orbital and spin cycles [i.e.\\ the frequency 2(\\beat) where $\\omega$ and $\\Omega$ are the spin and orbital frequencies respectively]. Note, though, that no other IP shows a lightcurve containing 2(\\beat) but not \\beat.  Other than this, \\vcen\\ is little studied, with, as yet, no ephemerides for the periodicities, no estimate of the field strength, and no determination of the binary inclination or of the component masses. Note that a possible grazing eclipse reported by Allan \\etal\\ (1999) was an artefact of incorrect data reduction. ",
        "conclusions": "We have presented phase-resolved spectroscopy of the intermediate polar \\vcen. We summarise here the strengths and weaknesses of the two models proposed for this system.  \\subsection{Partial-disc model} The strengths of the model are: (1) Disc-fed accretion explains an X-ray lightcurve varying only at the spin period. (2) The line-profile variations can be plausibly explained by invoking stream--disc overflow, in a manner seen in SW~Sex stars (although a weak X-ray beat pulsation might then be expected). (3) The beat-resolved tomograms can be interpreted as showing structure from the stream overflowing the disc, illuminated once per beat cycle by the X-ray beam.  The weaknesses are: (1) There is no explanation for the anomalously long spin period of \\vcen, except to claim that the system is not in equilibrium. (2) There is no easy explanation for the 1860-s optical pulsation.  \\subsection{Discless model} The strengths of the model are: (1) Explains the long spin period of \\vcen. (2) Explains the hook-like features in the tomograms, and the changes in the feature over beat-cycle phase (though there are still differences with the data). (3) Explains the 1860-s optical pulsation as ejection events occurring at 2(\\beat).  The main weakness is the fact that the X-ray lightcurve varies only at the spin frequency, and not at the orbital and beat frequencies, which argues against stream-fed accretion.  \\subsection{Further work} The summary above shows that we can't yet decide between the two models.  Note that there is also a scenario in which both may be `right': if \\vcen\\ had been discless for most of the past 10$^{6}$--10$^{7}$ y, this would explain the long spin period, even if a disc had formed more recently.  Observations needed to make further progress include a determination of the inclination of \\vcen, allowing us to match the observed line velocities to the models. Developments to the theory could include the addition of radiation processes to the diamagnetic-blob model, allowing a better comparison with line profiles, and computations of the predicted X-ray lightcurves. Also useful would be deeper searches for polarisation, following the initial work by Buckley \\etal\\ (1998).  In the diamagnetic-flow model the field is likely to be an order-of-magnitude stronger than in a disc-fed system of the same orbital period (fig 7 of King \\&\\ Wynn 1999), and further the polarised light would not be diluted by a bright disc, leading to a greater likelihood of detecting polarisation in these systems than in most IPs."
    },
    "0201/astro-ph0201197_arXiv.txt": {
        "abstract": "We investigate the possibility to find evidence for planets in circumstellar disks by infrared and submillimeter interferometry. We present simulations of a circumstellar disk around a solar-type star with an embedded planet of 1 Jupiter mass. The three-dimensional (3D) density structure of the disk results from hydrodynamical simulations. On the basis of 3D radiative transfer simulations, images of this system were calculated. The intensity maps provide the basis for the simulation of the interferometers VLTI (equipped with the mid-infrared instrument MIDI) and ALMA. While MIDI/VLTI will not provide the possibility to distinguish between disks with or without a gap on the basis of visibility measurements, ALMA will provide the necessary basis for a direct gap detection. ",
        "introduction": "Based on studies of the evolution of protoplanets in protoplanetary disks, it has been established that - depending on the hydrodynamic properties of the planet and the disk - giant protoplanets may open a gap and cause spiral density waves in the disk (see, e.g., Kley 1999, Kley et al.\\ 2001). The gap may extend up to several AU in width. Thus, the question arises if one can find such a gap as an indicator for the presence of a protoplanet with present-day or near-future techniques. In order to study this possibility, we use hydrodynamical simulations of a protoplanetary disk with an embedded planet (\\S~\\ref{model}) and compute the expected brightness distributions, with a 3D radiative transfer code (\\S~\\ref{rttech}). The resulting images are presented in \\S~\\ref{images}. Finally, we show that ALMA will provide the necessary basis to detect gaps in circumstellar disks in the mm/submm wavelength range (\\S~\\ref{simmial}). ",
        "conclusions": "\\label{concl}  Based on hydrodynamical simulations and subsequent 3D continuum RT calculations we generated millimetre images of a circumstellar disk with an embedded planet surrounding a solar-type star. The gap can be seen very clearly in the simulated images but an extremely high angular resolution is required ($\\approx$\\,10\\,mas). Because of the extreme brightness contrast in the innermost region of the disk in the near to mid-infrared, the gap can hardly be detected in this wavelength range with imaging/interferometric observations.  We found that it will be possible to distinguish between a disk with or without accretion onto the central star with MIDI, but the visibility at 10\\,$\\mu$m is almost insensitive to the presence of a gap. In contrast to this, the millimetre interferometer ALMA which will become available in the near future will achieve this goal: it will provide the basis for the reconstruction of an image of a gap. Thus, the search for massive protoplanets in circumstellar disks can be based on the indication of a gap. For this purpose, the longest possible baselines are required."
    },
    "0201/astro-ph0201368_arXiv.txt": {
        "abstract": "Spiral shocks in cataclysmic variables (CVs) are the result of tidal interactions of the mass donor star with the accretion disc. Their study is fundamental for our understanding of angular momentum transfer in discs. In our quest to learn how widespread amongst binaries spiral shocks are, and how their presence depends upon orbital period and mass ratio (as they are created by direct interaction with the donor star), we have obtained spectra of a large sample of CVs during their high-mass-transfer states. We find that 24 out of the 63 systems observed are candidates for containing spiral shocks. 5 out of those 24 CVs have been confirmed as showing shocks in the disc during outburst. ",
        "introduction": "In 1996, Sawada, Matsuda, \\& Hachisu suggested that the angular momentum in accretion discs could be transferred by means of shocks, tidally excited by the donor star. These shocks were seen for the first time in IP~Peg during outburst by Steeghs, Harlaftis, \\& Horne (1997). The shocks have since been seen in several other outbursts of IP~Peg (Harlaftis et al. 1999; Morales-Rueda, Marsh, \\& Billington 2000). The spiral pattern is not seen during quiescence (Marsh \\& Horne 1990), consistent with the idea that a hot disc is needed to lower the Mach number.  Spiral shocks have been successfully modelled using numerical calculations (Steeghs \\& Stehle 1999) but there are still some unanswered questions about these spiral patterns: How do they evolve with time? Do they form every outburst?  Are they still present on the return to quiescence? Above all we would like to know how widespread amongst other systems these shocks are, and how they depend upon the system's orbital period and mass ratio.  At present, IP~Peg provides the best example of these shocks, with some signs in the systems SS~Cyg (Steeghs et al.\\ 1996), EX~Dra (Joergens, Spruit, \\& Rutten 2000), U~Gem (Groot 2001) and WZ~Sge (Steeghs et al. this volume). Two common features seen in the spectra of these 5 systems during outburst are that they contain emission lines, and the shocks are stronger in He{\\sc ii}\\,4686\\AA. The small number of systems where shocks have been seen is not surprising as there are remarkably few extensive datasets of outburst spectra.  This work aims at solving this problem by obtaining spectra of dwarf novae during outburst and identifying those with strong emission lines, mainly He{\\sc ii}\\,4686\\AA. ",
        "conclusions": ""
    },
    "0201/astro-ph0201532_arXiv.txt": {
        "abstract": "{In this paper we report and discuss the results of a radio survey in the A3571 cluster complex, a structure located in the Shapley Concentration core, and formed by the three clusters A3571, A3572 and A3575. The survey was carried out simultaneously at 22 cm and 13 cm with the Australia Telescope Compact Array, and led to the detection of 124 radio sources at 22 cm. The radio source counts in this region are in agreement with the background counts. Among the 36 radio sources with optical counterpart, six have measured redshift which places them at the distance of the A3571 cluster complex, and nine radio sources have optical counterparts most likely members of this cluster complex. All radio galaxies emit at low power level, i.e. logP$_{22cm}$ (W Hz$^{-1}$) $\\le$ 22.6. A number of them are likely to be starburst galaxies. The radio luminosity function of early type galaxies is in agreement with that derived by Ledlow \\& Owen (1996) if we restrict our analysis to the A3571 cluster. On the basis of the multiwavelength properties of the A3571 cluster complex, we propose that it is a very advanced merger, and explain the radio properties derived from our study in the light of this hypothesis. ",
        "introduction": "Merging between clusters is thought to be responsible for significant changes in the physics of the intracluster medium and in the emission properties of the galaxy population, as a consequence of the enormous energy budget involved ($\\sim 10^{60}$ ergs).  The effects of this dynamical process on the hot gas are clear and well documented (Sarazin 2000), as well as the connection between cluster merging and  peculiar radio sources such as relics and halos (Feretti \\& Giovannini 1996). In particular, these extended radio sources are found to be associated with clusters with some degree of disturbance, suggesting that merging is responsible for the reacceleration of the relativistic electrons (Feretti \\& Giovannini, 2001; Ensslin\\& Br\\\"uggen, 2001; Brunetti et al. 2001; Buote 2001). On the other hand it is not yet understood whether cluster--cluster collision is able to modify the radio emission properties of a single galaxy and/or affect the statistics of the radio source population.  In order to address this issue, we are carrying on a multiwavelength study of the central part of the Shapley Concentration, where an unusually high number of clusters are merging. This situation originated from the high local density, that induced peculiar velocities of the order of $\\sim 1000$ km s$^{-1}$.  Zucca et al. (1993) performed a percolation analysis of the Abell-ACO clusters in order to find superclusters\tand to study the grouping characteristics as a function of the overdensity with respect to the mean density of clusters. In particular, they found that at the highest density contrasts ($>40$), the central part of the Shapley Concentration is fragmented in three main structures  (``cluster complexes\" or ``groups of clusters\"). One is dominated by the cluster A3558 and could be considered the core of the supercluster; a second is formed by A3528, A3530 and A3532. Both complexes present very clear signs of ongoing merging (Hanami et al. 1999; Bardelli et al. 1996; Ettori et al. 1997; Kull \\&B\\\"ohringer 1999; Schindler 1996). The third group of clusters in this region is dominated by A3571, and includes A3572 and A3575.  We surveyed the A3558 and A3528 cluster regions at 22 cm with the Australia Telescope Compact Array (ATCA) to the limiting flux density of $\\sim 0.8$ mJy (Venturi et al. 2000 and 2001). It was found that the A3558 complex is characterised by a significant lack of radio sources with respect to the cluster sample of Ledlow \\& Owen (1996), suggesting the possibility that cluster--cluster collisions would ``swith off\"the existing radio sources, or inhibit the birth of new ones. On the contrary, the number of radio sources found in the A3528 complex is statistically consistent with the one of the control sample. These two different behaviours can be possibly explained in terms of merging age: the A3558 complex is at an advanced stage (possibly after the first core-core encounter, Bardelli et al. 1998), while the A3528 complex is an early merging, where clusters just started to interact (see for example the simulations by Ricker \\& Sarazin 2001).  In this paper we present the radio--optical analysis of the third group of clusters in the central part of the Shapley Concentration, i.e. the A3571--A3572--A3575 chain. In Section 2 we overview optical and X--ray properties of the A3571 complex; in Section 3 we report on the radio observations; the radio source sample and the optical identifications are described in Section 4; the radio properties of the A3571 complex are presented in Section 5, and the radio and optical number counts are derived in Section 6. Discussion and conclusions are given in Section 7.  We will use H$_o$ = 100 km s$^{-1}$Mpc$^{-1}$ and q$_o$ = 0, which leads to a linear scale 1 arcsec = 0.55 kpc at the average redshift of the cluster complex (z = 0.039).  \\medskip ",
        "conclusions": "This work is part of a larger project whose aim is to study the effects of cluster mergers at radio wavelengths in the central region of the Shapley Concentration. The results obtained for the two major merging complexes have already been presented and discussed (Venturi et al. 2000 and 2001). Here we report on ATCA 22 cm observations of the third merging complex in the Shapley Concentration, formed by the three clusters A3571, A3572 and A3575. In particular in this paper we have focussed on the possible influence of cluster mergers on the statistical properties of the radio galaxies. We can briefly summarise our results as follows.  \\medskip \\noindent {\\it (a)} Six radio galaxies from our 22 cm radio sample were identified with galaxies in the in the A3571 chain. Other nine, associated with galaxies brighter than b$_j \\le 18.5$ and with no measured redshift, are possible Shapley member candidates.  \\noindent {\\it (b)} All Shapley radio galaxies and candidates are weak radio sources, with radio powers in the range $21.2 < $ logP$_{22}$ (W Hz$^{-1}$) $ < 22.6$. They are located in A3571 and in A3575 (none in A3572).  \\noindent {\\it (c)} The estimate of the RLF based on the analysis by LO96 points towards a lack of radio galaxies in the A3571 complex when we normalise the number of radio galaxies to the whole chain. However, if we take into account that all the radio galaxies considered in the RLF are located within 0.5 R$_A$ from the centre of A3571, the number of detections is in very good agreement with the expectations.  \\noindent {\\it (d)} The radio source counts in this region are dominated by the background counts, despite an optical overdensity of $\\sim 1.5$.   \\medskip Our results suggest a dual character of the A3571 cluster complex, where the A3571 cluster alone shows different properties than the chain as a whole, as clear from points {\\it (b)} and {\\it (c)} above. The same dual properties are evident also from observations at other wavelengths.  The cluster gas in A3571 is very hot (T =7.6 keV) and luminous (L$_{X,bol}~\\sim~4\\times 10^{44}$ erg s$^{-1}$, Ettori et al. 1997). With the observed temperature we can estimate the expected galaxy velocity dispersion, according to the following formula (Lubin \\& Bahcall, 1993):\\\\   \\hspace{-5mm}$ \\sigma_v = 10^{2.52\\pm0.07}(kT)^{0.6\\pm0.11}$\\\\  \\noindent which gives a value $\\sigma_v \\sim 1118$ km s$^{-1}$. This is in good agreement with $\\sigma_v \\sim 1022$ km s$^{-1}$, obtained by Quintana \\& De Souza (1993) on the basis of redshift measurements, and suggests that A3571 is relaxed, as proposed by Nevalainen et al. (2000). Another piece of evidence in favour of virialisation is the presence of a cooling flow in its innermost region (Peres et al. 2000), which requires that the cluster has been in equilibrium at least since 4--6 $\\times 10^9$ yrs.  Conversely, the redshift survey of Quintana et al. (1997), together with the presence of the giant cD galaxy at the centre of A3571 (MGC05$-$33$-$002) suggest that the whole chain is not relaxed and that a recent merger may be responsible for the formation of the central cD galaxy.  Inspection of the optical isodensity contours given in Fig. \\ref{contour} shows a major optical overdensity in A3571, while the other two clusters are not as well defined (especially A3575). In the light of the optical distribution and of the results summarised above, it is very important to understand whether A3575 and A3572 are the debris of a merger event, or if they are smaller entities yet to interact with the massive A3571.  \\medskip The results obtained from our 22 cm survey may help in clarifying the situation.  It has been argued that a connection exists between cluster mergers and a high number of starburst radio galaxies. In particular Owen et al. (1999), in a comparative study of the two clusters A2125 and A2645, suggested that merging could be responsible both for the high fraction of blue galaxies and of radio AGN in A2125, compared to the relaxed A2645.  The results presented in this paper are not conclusive in this respect, mainly because of the lack of complete spectroscopy in this region of the sky, however there is indication that a considerable fraction of candidate starburst galaxies is located in A3571. In particular, we found 9 galaxies brighter than b$_j \\le 18.5$ (including confirmed members and candidates) characterised by low radio power, i.e. logP$_{22}$ (W Hz$^{-1}$) $\\le 21.78$, whose radio emission could be driven by star formation. These numbers should be compared to the 12 - 13 starburst galaxies detected in the richness 4 distant cluster A2125 by Owen et al. (1999). It is crucial that redshift and spectral information becomes available for the objects in the A3571 complex, to confirm that we are indeed observing a high fraction of starburst galaxies. \\noindent We note that this result is considerably different from what we obtained for the A3528 and A3558 complexes, where no hint of such excess was revealed (Venturi et al. 2000 and 2001) at similar radio power limits.  We propose that A3571 is the result of a recent merger event, at the very last stage of its evolution, and that the three main cluster complexes in the central region of the Shapley Concentration are part of an evolutionary sequence. In particular, the wealth of information available from radio to optical (both photometry and spectroscopy), up to X--ray energies, suggest the following scenario:  \\noindent {\\it (i)} the A3528 cluster complex is at the very beginning of a merger event, where the two merging entities have just started ``to feel each other''. The gradients in the temperature distribution of the intracluster gas delineate the merging front (Schindler 1996). The radio luminosity function of elliptical galaxies is in good agreement with that of ellipticals not located in merging clusters (LO96), and no sign of starburst emission, possibly induced by merger shocks, is detected (Venturi et al. 2001). We suggested that the pre-merging stage had not yet had time to affect the radio emission properties of the cluster galaxies in the complex.   \\noindent {\\it (ii)} the A3558 complex is thought to be an advanced merger, where two clusters of similar mass have already undergone the first core--core encounter. The amazing similarity between the galaxy distribution (Bardelli et al. 1998) and the gas density distribution (Ettori et al. 1997; Kull \\& B\\\"ohringer 1999) provides further evidence of the ongoing process. In the radio band it was found (Venturi et al. 2000) that this complex shows a significant deficit of radio galaxies in comparison with the radio luminosity function of LO96, suggesting that the major encounter switched off the nuclear radio emission and temporarily inhibited its formation. No radio excess of starburst origin was detected in the shock region, however data from a deeper survey in the same region are being analysed (Venturi et al. in prep.).  \\noindent {\\it (iii)} We suggest that the A3571 complex represents the final stage of a merger event, where A3571 itself is the resulting cluster after virialization of the merger. The distribution of the gas is already relaxed, as well as the galaxy distribution in A3571, while the outer edge of the galaxy distribution is still unrelaxed. The radio properties reflect the different dynamical stage of the central relaxed region of the complex (the cluster A3571) and the active dynamics of the outskirts. The location of the radio galaxies in A3571 suggests that they had time to develop a nuclear source after the active merging ceased.  We note that the distribution of the optical galaxies and of the X--ray emitting gas in these three cluster complexes is remarkably similar to the various stages of the cluster--cluster collision recently modelled by Ricker \\& Sarazin (2001). Their simulations for a frontal merger of two clusters with mass ratio M$_1$/M$_2$ = 1 show an A3528--like situation for t=0, which evolves into an A3558 scenario at t $\\sim$ 5 Gyr, to end up in the A3571 case after $\\sim$ 9 Gyr."
    },
    "0201/astro-ph0201538_arXiv.txt": {
        "abstract": "Although a large number of orbital periods of cataclysmic variable stars (CVs) have been measured, comparison of period and luminosity distributions with evolutionary theory is affected by strong selection effects. A test has been discovered which is independent of these selection effects and is based upon the kinematics of CVs (Kolb \\& Stehle, 1996). If the standard models of evolution are correct then long-period ($P_{\\rm orb} > 5$hrs) CVs should be typically less than 1.5 Gyr old, and their line-of-sight velocity dispersion ($\\sigma_\\gamma$) should be small. We present results from a pilot study which indicate that this postulate is indeed true. Four long-period dwarf novae (EM~Cyg, V426~Oph, SS~Cyg and AH~Her) were observed over a complete orbit, in order that accurate radial velocities be obtained. We find values of -1.7, 5.4, 15.4 and 1.8 \\kms\\ with uncertainties of order $3$\\,\\kms, referred to the dynamical Local Standard of Rest (LSR), leading to a dispersion of $\\sim$8\\kms. Calculation of a 95 per cent confidence interval gives the result $4 < \\sigma_\\gamma < 28$\\kms\\ compared to a prediction of $15$\\,\\kms. We also have an improved determination of mass~donor spectral type, $K_{2}$ and $q$ for the four systems. ",
        "introduction": "Dwarf novae (DN) are a sub-group of cataclysmic variable stars. They consist of a white dwarf (primary) star accreting matter from a red dwarf (hereafter `mass~donor') which is in contact with its Roche lobe. The material streams from the mass~donor through the inner Lagrangian point into the potential well of the primary, where it impacts onto the outer edge of an accretion~disc. Observable features which distinguish DN from other CVs are the outbursts which occur on a regular basis (with a baseline from days to months) and cause the system to increase in brightness by $m_{v} \\approx 2 - 5$ magnitudes.  The mass transfer process in CVs is driven by orbital angular momentum loss. Gravitational radiation is highly effective at short orbital periods, and theoretical calculations reproduce well the observed mass transfer rates in CVs. For longer period systems, with higher mass transfer rates, an alternative mechanism is needed. This is widely accepted to be {\\it magnetic braking}, a spin-down effect on the mass~donor due to a stellar wind that effectively co-rotates with its magnetic field. Due to effective tidal locking in CVs, the mass~donor will not slow down, and instead angular momentum is extracted from the binary orbit. However, magnetic braking still remains only a hypothesis, and the level of direct observational support is still low.  In the standard model for the formation and evolution of CVs magnetic braking is only active for donor stars with a radiative core. Hence there is a marked difference between the evolutionary time scale of systems above the period~gap ($P_{\\rm orb} \\ga 3$\\,h) and those below the gap ($P_{\\rm orb} \\la 2$\\,h). Kolb \\& Stehle~\\shortcite{kolste} used population synthesis methods to confirm this. They determined the age structure of a model population of Galactic CVs, and by convolving this with the observed age-space velocity relation of Wielen~et~al.~\\shortcite{wie} (which is believed to result from the diffusion of stellar orbits due to gravitational interactions with massive objects) obtained the theoretical distribution of systemic, hereafter $\\gamma$, velocities. They showed that the age of a system in the present CV population is largely determined by the time-scales of the orbital angular-momentum loss mechanisms. In addition, the models showed that the brightness-dependent selection effects which have hitherto plagued comparisons between observations and theory preserve the age differences, thereby providing an opportunity to test directly the magnetic braking model.  So, if the standard models are correct, then the CVs having periods longer than the upper limit of the period~gap ($P_{\\rm orb} \\ge 3$\\,hours) should be younger ($\\le 1.5$Gyr), and therefore have a smaller line-of-sight velocity dispersion according to the empirical age--velocity dispersion relation (predicted value $\\sim 15$\\,\\kms).  Conversely, those CVs with orbital periods shorter than the lower limit of the period~gap, should be older ($\\ge$ 3-4Gyr) and show a larger velocity dispersion (predicted as $\\sim 30$\\,\\kms). \\begin{figure} \\hspace*{\\fill} \\psfig{file=figure1.ps,width=85mm} \\hspace*{\\fill} \\caption{The figure (from Kolb \\& Stehle, 1996) shows the relative number of systems as a function of age, for CVs above the period gap ($P_{\\rm orb}>3$hr, dotted line) and below the gap ($P_{\\rm orb} < 2$ hr, dashed line). The solid line is the total fraction of CVs above the period gap as a function of age. The relative youth of the longer period systems is a firm prediction of the magnetic braking model of CV evolution.} \\label{fig:ages} \\end{figure}  In a review by van Paradijs, Augusteijn \\& Stehle~\\shortcite{vanP}, the observed $\\gamma$ velocities for a sample of CVs were collected from published radial-velocity studies and statistically analysed. They could detect no difference between the velocity dispersions for below-gap and above-gap objects~(see section~\\ref{sec-dis}). All their data were taken from the literature, and they assumed that there was no significant difference between measurements taken from emission and absorption line measurements in their sample. The problems associated with using values taken from the literature are not negligible. The $\\gamma$ velocity is a measure of the centre-of-mass radial velocity of the binary star, and can be deduced directly from radial velocity curves.  Typically with the emission lines from CVs, a $\\gamma$-value which is only one tenth the size of the radial velocity variations in the system, and only one hundredth of the total spectral line width, is being measured.  The radial velocities obtained from emission lines are well known to be affected by the internal motions in the disc or stream. Therefore, the measurements of $\\gamma$ obtained from emission lines may not accurately reflect the motion of the white dwarf at all. This will be especially true if the inner disc is distorted in some way, for example, in magnetic CVs where the field on the primary is so strong it disrupts the accretion disc and disrupts the accretion disc near to the surface of the white dwarf. Measuring radial velocities from the absorption lines due to the mass~donor helps reduce potential sources of error. However, these absorption lines are not visible in all dwarf~novae. They are generally only present in those systems with orbital periods above the period~gap.  Many radial-velocity curves for % different CV systems have been constructed. However, if the only requirement is a measurement of the radial velocity semi-amplitude, care will not have been taken to ensure that systematic errors do not dominate. It is vital that spectra are adequately sampled, and that a reasonable number of radial velocity standards have been observed. If a lower resolution is used to obtain data, then the probability of the spectral lines being affected by blending is significantly increased.           Here, we present the initial results of a project to deduce reliable $\\gamma$ velocities for a large sample of longer-period ($P_{\\rm orb} \\ge 6$ hours) non-magnetic CVs. The aim of this initial investigation was to test the methods and procedures described in \\S III, to see whether they could produce sufficiently accurate absolute $\\gamma$ velocity values, capable of being used to construct an observed velocity-dispersion relation for Galactic CVs, and thus for comparing directly with the theory. ",
        "conclusions": "Four long period ($P_{\\rm orb} > 5$ hours) dwarf~novae have been observed over a complete orbital period, in order to determine accurate absolute systemic velocities (to within $\\pm5$\\,\\kms). On this initial sample, the aim was to test thoroughly the methods being used to obtain the velocities, to make sure that any possible sources of error were effectively minimised. Sources of systematic error are fairly numerous using these methods, however they are very easy to identify and minimise, and in certain cases to eliminate altogether. A large sample of radial-velocity standard stars were observed, to ensure that so-called `template mismatch' did not introduce any significant error. This enabled us to compile a list of reliable standard stars to use for future observing runs for this project. In addition, because of the high accuracy requirement for measurement of the systemic velocities, accurate measurements of other binary parameters can be made. For example, high quality radial-velocity curves mean that good semi-amplitude measurements can be obtained. Existing orbital periods can also be checked, and a value for the projected rotational velocity of the mass~donor can be determined, implying that a value for the mass ratio, $q$, can be calculated. This value of $q$ is then independent of the semi-amplitude value determined from the emission lines (measurements of which are uncertain and prone to error), and can then be used to determine the component masses. Doppler maps have been constructed from the H$\\alpha$ line profiles obtained in this study. Due to the coverage of a full orbital period, an orbit-averaged Doppler image has be created for each object. These have proved surprising.  Further evidence for the existence of ``sub-Keplerian'' velocities is seen in the maps, as are peculiar features which may be be attributable to an outflow from the systems. These features may be peculiar to H$\\alpha$ emission; further investigation is needed.  Our initial results suggest that the observations agree with the theory (as set out by Kolb \\& Stehle, 1996). The observed velocity dispersion for the sample is $\\sim 8 \\kms$, which is smaller than the predicted value of $15 \\kms$. If, with further observations, this value does not alter significantly, then we can infer that the longer period systems are a young subgroup, with an average predicted age of $\\le 1.5 {\\rm Gyrs}$. Once systemic velocities for the longer-period systems have been adequately constrained then our understanding of their kinematics and evolution will be more complete, and we can begin the search for the predicted significant difference between the velocity dispersions of CVs below and above the period gap."
    },
    "0201/astro-ph0201012_arXiv.txt": {
        "abstract": "{ Scattering and absorption of light by a homogeneous distribution of intergalactic large dust grains has been proposed as an alternative, non-cosmological explanation for the faintness of Type Ia supernovae at $z\\sim 0.5$. We investigate the differential extinction for high-redshift sources caused by extragalactic dust along the line of sight.  Future observations of Type Ia supernovae up to $z\\sim 2$, e.g. by the proposed SNAP satellite, will allow the measurement of the properties of dust over cosmological distances. We show that 1\\% {\\em relative} spectrophotometric accuracy (or broadband photometry) in the wavelength interval 0.7--1.5 $\\mu$m is required to measure the extinction caused by ``grey'' dust down to $\\delta m=0.02$ magnitudes.  We also argue that the presence of grey dust is not necessarily inconsistent with the recent measurement of the brightness of a supernova at $z=1.7$ (SN 1997ff), in the absence of  accurate spectrophotometric information of the supernova. ",
        "introduction": "There is observational evidence (\\cite{scp,highz}) that Type Ia supernovae, when used as calibrated standard candles, are dimmer at high redshift than can be explained in models without a cosmological constant.  At least there seems to be a need for a non-clustered ``dark energy'' component with a negative coefficient in the equation of state, such as obtained in models with an evolving scalar field (``quintessence'') (\\cite{steinhardt:haga}). Since there are potential systematic effects affecting this interpretation, it is important to investigate alternative explanations. In this note, we investigate dimming due to absorption and scattering on intergalactic dust, as has been proposed by Aguirre (\\cite{aguirre99a,aguirre99b}) to be a viable explanation for the supernova results. The recent measurements of the CMB small angle anisotropies showing that the universe is most likely flat (\\cite{boomerang,maxima,dasi}), combined with measurements of $\\Omega_{\\rm M}\\sim 0.3$ from large scale structure (\\cite{2DF}) and galaxy cluster evolution (\\cite{bahcall}) makes Aguirre's idea for the origin of the Type Ia supernovae faintness at $z\\sim 0.5$ unlikely.  Moreover, \\cite{aguirrehaiman} have shown that the amount of dust required to make the supernova results compatible with a flat universe, as indicated by the CMB results, but {\\em without} a cosmological constant, is already disfavored by the far-infrared background measured by the DIRBE/FIRAS instruments. However, a sizable grey dust column density capable of biasing the results cannot be excluded with the present knowledge. Precision measurements of ``Dark Energy'' and ``Dark Matter'' with high-z supernovae have been proposed see e.g. (\\cite{goliath,mortsell}). The assumption in those studies is that the  systematic uncertainties do not exceed $\\delta m$=0.02 mag. Thus, in this work we concentrate on the needed relative spectrophotometric accuracy to meet this requirement.   Extinction  must be considered  for at least four different dust environments: \\begin{enumerate} \\item A homogeneous intergalactic dust component. Particular attention must be paid to the possibility that the mechanisms expelling dust from galaxies and clusters destroy the smallest grains, thereby causing very little reddening (``grey dust''). \\item A host galaxy dust component. \\item Dust in galaxies between the source and the observer. \\item Milky-Way dust \\end{enumerate}  Properties such as the extinction scale-length and the wavelength dependence can be different for each case.  The main emphasis of this work is on the calculation of the possible effects from intergalactic dust (1). To avoid observational constraints on reddening, such a component must mainly consist of large dust grains as described in  (Aguirre \\cite{aguirre99a,aguirre99b}). However, even large-grain dust will cause {\\em some} reddening, and going to higher redshift should enable to distinguish between extinction or a cosmological origin  for the faintness of  Type Ia supernovae at z$\\sim$0.5. The future data sample we have in mind here is the one expected from the SNAP satellite (\\cite{snap}), which will provide several thousand Type Ia supernovae out to $z\\sim 1.7$.  In this note we also discuss the light extinction by ``normal'' dust in either the SN host galaxy (2) or galaxies along the line of sight (3). As we will see, the latter only affects sizably about 1 \\% of the supernovae at $z\\sim 1$. ",
        "conclusions": "The effects of grey dust extinction capable of biasing the results of experiments such as the proposed SNAP satellite can be diagnosed through accurate relative spectrophotometry or broadband photometry at the 1\\% level in the $0.7-1.5$ $\\mu$m wavelength range, as shown in e.g. Figs.\\ \\ref{fig:omp3oxp7rv9p5_2} and \\ref{fig:color_omp3oxp7rv9p5_2}. This would allow tests for intergalactic dust obscuration affecting the measurement of high-z supernovae up to $\\Delta m \\sim 0.02$, the target for systematic uncertainties for SNAP.  Achieving 1-2\\% relative spectrophotometric accuracy in the optical and NIR for 22-29 magnitude objects is rather challenging. As the measurements will rely on a large number of homogeneous high-redshift sources, it is not required to get S/N=50-100 for individual objects. On the other hand, the instrument calibration must be within 1\\% over the course of at least one year for the case of SNAP.  This can be accomplished e.g. through repeated observations of hot galactic white dwarfs (\\cite{bohlin}; \\cite{finley}). In the Rayleigh-Jeans limit, T$_{\\rm eff}\\gg$20000 K, the color of the calibrators become independent of temperature and are therefore ideal for relative calibration.  Significant progress in examining the possible bias due to grey dust in the published Type Ia supernova Hubble diagram can already be made with existing NIR instruments. If the faintness of Type Ia SNe at high-z is to be attributed to grey dust obscuration as opposed to the cosmological explanation a color extinction $E(R-J)^{\\rm obs} \\gsim 0.1$ mag is to be expected for sources at $z\\sim0.5$. Testing this possibility is within reach with ground based facilities.  We have also examined the possibility of extinction as the beam of high redshift passes through foreground galaxies. This was found to be a relatively small effect, only causing $>0.02$ mag dimming for less than 1 \\% of the sources at $z=1$, whereas extinction in the host galaxy potentially is a more serious problem.  However, in general we expect dust in galaxies along the line of sight and in host galaxies to have $R_V\\sim 3$, causing considerably more reddening than an intergalactic \"grey\" dust component. Thus, with high accuracy spectrophotometry, it should be possible to control the effects from extinction in galaxies."
    },
    "0201/astro-ph0201362_arXiv.txt": {
        "abstract": "Recent work has shown that it is possible to systematize quasar spectral diversity in a parameter space called  ``Eigenvector 1'' (E1).  We present median AGN spectra for fixed regions of the E1 (optical) parameter space (FWHM(\\hb) vs. equivalent width ratio \\rfe=W(\\feiiq)/W(\\hb). Comparison of the median spectra for different regions show considerable differences. We suggest that an E1-driven approach to median/average spectra emphasizes significant differences between AGN, and offers more insights into AGN physics and dynamics than a single population median/average derived from a large and heterogeneous sample of sources. We find that the \\hb\\ broad component line profile changes in shape along the E1 sequence both in average centroid shift and asymmetry. While objects with FWHM(\\hbbc)$\\la$ 4000 \\kms\\ are well fitted by a Lorentz function, AGN with FWHM(\\hbbc)$\\ga$ 4000 \\kms\\ are well fitted if two broad line components are used: a broad (the ``classical'' broad line component) and a very broad/redshifted component. ",
        "introduction": "Studies of broad emission line spectra for AGN provide both the strongest constraint on models for the nebular physics and kinematic models of the broad line emitting clouds.  Ideally one would like to have both high signal-to-noise ratio (S/N) and moderate resolution (\\ltsima 5 \\AA) measures for the strongest high and low ionization lines as well as measures across as wide a wavelength range as possible in order to better characterize the continuum shape.  The former measures have led to the Eigenvector 1 parameter space concept (Sulentic et al.  2000a,b) that is, in part, built upon foundations laid almost ten years ago (Boroson \\& Green 1992; hereafter BG92). The above goals are very observing time intensive but recent AGN surveys have provided an alternate method for examining quasar spectra from 1000 -- 10000 \\AA. Average or composite spectra derived from large survey databases involving hundreds of AGN (LBQS: Francis et al. 1991; HST: Zheng et al. 1997; 1998; FIRST:  Brotherton et al. 2001; Sloan survey: Vanden Berk et al. 2001) make it possible to generate a ``typical'' quasar spectrum from blueward of Ly$\\alpha$\\ to \\ha.  What is unclear is whether such composite spectra are astrophysically useful beyond, perhaps, allowing us to identify lines that are too weak to be seen in individual source spectra.  The fundamental question is whether the similarities or the differences in AGN line and continuum phenomenology tell us more about the underlying physics. The Eigenvector 1 concept has been advanced as a possible ``H-R Diagram\" for AGN in the sense that it appears to provide parameter space discrimination between all major classes of broad line sources as well as a correlation for, at least, radio quiet sources (Sulentic et al. 2000b).  The correlation and distribution of radio-quiet sources (with FWHM(\\hb)$\\leq$ 4000 \\kms) in E1 have been reasonably well fit with a model that sees accretion rate, convolved with source orientation, as the principal physical driver (Marziani et al.  2001). The generality of E1, of course, requires much further testing but unless it is far off the mark, it suggests that average or composite spectra should be viewed in the same way that one would view an average stellar spectrum taken over the full effective temperature range that is observed in the H-R Diagram (spectral types O-M). Interpretation of composite spectra from heterogeneous samples will be complicated in two ways: (1) they will be  subject to selection biases dependent on the relative number of sources with each ``spectra type'' in a given sample and, more significantly, (2) they will average the spectra of sources with dramatically different physical properties.  We suggest that the most useful approach to averaging AGN spectra lies within the E1 context. We present average spectra for fixed domain quadrants in E1 followed by a brief discussion of important differences. The emphasis is on the FWHM(\\hb) measures although large differences in the equivalent width parameter are also seen. Full discussion of the equivalent width differences must wait for significant numbers of very high S/N measures for \\feii\\ weak sources. ",
        "conclusions": "We generated average QSO spectra  in the 4200--5700 \\AA\\ spectral region for fixed optical parameter bins in E1. In fixed parameter range bins the mean spectral type appears to change systematically across E1.  This systematic behaviour is reflected in  E1 X-ray and UV measures as well. Composite spectra in an E1 context should provide useful input for theoretical modeling. Ignoring the diversity of the AGN \\hbbc\\ profile may make any theoretical modeling of the BLR kinematics and dynamics unreliable."
    },
    "0201/astro-ph0201154_arXiv.txt": {
        "abstract": "Many workers have found that the recollapse of a dark matter halo after decoupling has a self-similar dynamical phase. This behaviour is maintained strictly so long as the infall continues but it appears to evolve smoothly into the virialized steady state and to transmit some of its properties intact. The density profiles established in this phase are all close to the isothermal inverse square law however, which is steeper than the predictions of some n-body simulations for the central regions of the halo, which are in turn steeper than the density profiles observed in the central regions of some galaxies; particularly dwarfs and low surface brightness galaxies. The outer regions of galaxies both as observed and as simulated have density profiles steeper than the self-similar profile. Nevertheless there appears to be an intermediate region in most galaxies in which the inverse square behaviour is a good description. The outer deviations can be explained plausibly in terms of the transition from a self-gravitating extended halo to a Keplerian flow onto a dominant central mass (the isothermal distribution can not be complete), but the inner deviations are more problematic. Rather than attack this question directly, we use in this paper a novel coarse-graining technique combined with a shell code to establish  both the distribution function associated with the self-similar density profile and the nature of the possible deviations in the central regions. In spherical symmetry we find that both in the case of purely radial orbits and in the case of orbits with non-zero angular momentum the self-similar density profile should flatten progressively near the centre of the system. The NFW limit of $-1$ seems possible. In a section aimed at demonstrating our technique for a spherically symmetric steady state, we argue that a Gaussian distribution function is the best approximation near the centre of the system. ",
        "introduction": "\\label{sec:intro} The study of the radial infall of dark matter and the realization that the evolution becomes self-similar has by now a long history (see for example \\cite{HW99} for a summary of the relevant references and a nearly state of the art statement of our understanding ). The salient theoretical behaviour is well known if not completely understood, but it yields density profiles that are theoretically distinct from the inverse square law but observationally indistinguishable from this law. However the NFW law from n-body simulations \\cite{NFW} even if slightly steeper when resolution effects are taken into account (e.g. \\cite{Moore98}; but see also \\cite{Kravtsov98}) is distinctly flatter than the inverse square law in the central regions. Moreover in some galaxies (low surface brightness or LSB galaxies, DE Blok \\etal 2001; some dwarf ellipticals, Stil 1999; and even brighter elliptical galaxies compared to the fainter objects, Merritt and Cruz 2001) the observed density profile is significantly flatter in these regions than even the n-body profiles (except \\cite{Kravtsov98}). Both observed and simulated profiles are steeper than the inverse square law in the outer regions, but this can be understood \\cite{HW99} in terms of secondary accretion after most of the halo mass has fallen in.  We take the view in this paper that if finite resolution effects are at work in the centre of the simulations, and perhaps also in reality since obviously real particles can not yield an infinite density cusp; then we may gain some insight by solving analytically the Collisionless Boltzmann Equation (CBE) for the distribution function (DF) in an series expansion in powers of the inverse `smoothing length' (to be defined precisely below). In such an expansion the lowest order is the coarsest grained approximation and higher orders yield progressively finer grained information. It transpires that in this way we can deduce the radial distribution function necessary to maintain the strict self-similarity and we may place some constraints on modifications to this DF that can produce flattening.  Our method of coarse graining is by non-canonical transformation on the phase space (we make a special choice of a stretching transformation here, but the method is more general) which produces a definite equation to be satisfied by each term in our series. We illustrate the method first by applying a stretching transformation to a spherically symmetric isotropic equilibrium. The bulk of the paper is devoted to exploring the self-similar `equilibrium'maintained by the continuing accretion onto a relaxed core \\cite{FG84}; \\cite{B85}; \\cite{HW99}.  A consideration of the entropy for the system allows us to suggest that the central DF in spherical equilibrium is Gaussian in velocity with a separable radial dependence, while in exact self-similar secondary infall the DF is everywhere proportional to the square root of the particle binding energy. Moreover the requirement that the coarse grained entropy be maximal predicts perturbations that flatten the inner regions of steep self-similar profiles and steepen the outer regions of flat  self-similar profiles. We show also that the addition of angular momentum to spherically symmetric orbits does not change these conclusions.  \\renewcommand{\\textfraction}{0} \\renewcommand{\\topfraction}{1} \\renewcommand{\\bottomfraction}{1} ",
        "conclusions": ""
    },
    "0201/astro-ph0201318_arXiv.txt": {
        "abstract": "We study the evolution of disk accretion during the merger of supermassive black hole binaries in galactic nuclei. In hierarchical galaxy formation models, the most common binaries are likely to arise from minor galactic mergers, and have unequal mass black holes. Once such a binary becomes embedded in an accretion disk at a separation $a \\sim 0.1 \\ {\\rm pc}$, the merger proceeds in two distinct phases. During the first phase, the loss of orbital angular momentum to the gaseous disk shrinks the binary on a timescale of $\\sim 10^7 \\ {\\rm yr}$. The accretion rate onto the primary black hole is not increased, and can be substantially reduced, during this disk-driven migration. At smaller separations, gravitational radiation becomes the dominant angular momentum loss process, and any gas trapped inside the orbit of the secondary is driven inwards by the inspiralling black hole. The implied accretion rate just prior to coalescence exceeds the Eddington limit, so the final merger is likely to occur within a common envelope formed from the disrupted inner disk, and be accompanied by high velocity ($\\sim 10^4 \\ {\\rm kms}^{-1}$) outflows. ",
        "introduction": "A large fraction of galaxies in the local Universe harbor supermassive black holes \\citep{magorrian98}, with the black hole mass correlating well with the velocity dispersion of the galaxy \\citep{ferrarese2000,gebhardt2000}. One explanation for this correlation is that black holes, like galaxies themselves, have grown in mass through the hierarchical merging of smaller progenitors \\citep{haehnelt2000,cv2000,menou2001}. This implies that the frequent galactic mergers at high redshift could also involve the inspiral, and probably eventual coalescence, of supermassive black holes. Typically, these binaries will pair black holes of disparate masses, since minor galactic mergers are more common than major mergers in hierarchical models of galaxy formation \\citep{haehnelt2000}.  For a black hole binary with masses of $10^8 \\ M_\\odot$ and $10^6 \\ M_\\odot$, gravitational radiation will lead to merger within a Hubble time if the separation is $\\lesssim 10^{-2} \\ {\\rm pc}$. Reaching such a small separation, following a galactic merger, requires that the black holes lose almost all of their orbital angular momentum to stars, dark matter, or gas. Angular momentum loss to stars occurs readily at large radii, but slows down at pc-scale separations if the supply of stars on low angular momentum orbits, which can interact with the binary, becomes depleted \\citep{begelman80,mikkola92,makino97,quinlan97}. The extent of this loss cone depletion remains a subject of active debate \\citep{milo2001,yu2001,zhao01}, with recent work suggesting that stellar dynamical processes may suffice to bring binaries into the gravitational radiation dominated regime. Irrespective of these uncertainties, gas, if it is present at small radii in the nucleus, is likely to be the most efficient sink of binary angular momentum \\citep{begelman80,pringle91,ivanov99,gould2000}. In this {\\em letter}, we study the role of gas in hastening the merger process, and the influence of the binary on the accretion rate and luminosity of the black holes. ",
        "conclusions": "The ubiquity of supermassive black holes in the local Universe \\citep{magorrian98} suggests that black holes may also be common in the high redshift progenitors that merged to make up present day galaxies. Although a large black hole fraction at high-$z$ is not proven \\citep{menou2001}, a high frequency of mergers between galaxies harboring black holes would have interesting implications for the accretion history of AGN \\citep{haehnelt1998,haehnelt2000,cv2000}, and for the structure of galaxies themselves \\citep{ravindranath2001}.  In this {\\em letter}, we examined the coupled evolution of a supermassive black hole binary embedded in a gaseous accretion disk. First, we showed that, in a gas-rich environment, a binary with a separation of $0.1 \\ {\\rm pc}$ merges within around $10^7 \\ {\\rm yr}$. Close binary quasars are therefore expected to be uncommon. Long lived black hole binaries are more likely to be found in apparently quiescent galactic nuclei at low redshift, where the disk accretion rate is lower, and where migration may be further slowed by inefficient disk angular momentum transport \\citep{menou01}. Second, we examined the possible observational signatures of the merger. We argued that large enhancements to the accretion rate were unlikely during the phase of disk-driven migration. Forced accretion of any inner disk during the final, gravitational wave driven inspiral, however, is likely to lead to much of the inner disk being expelled in a high velocity quasi-spherical outflow."
    },
    "0201/astro-ph0201404_arXiv.txt": {
        "abstract": "In this paper, the influence of binary stars on the measured kinematics of dwarf galaxies is investigated. Using realistic distributions of the orbital parameters (semi-major axis, eccentricity, \\ldots), analytical expressions are derived for the changes induced by the presence of binary stars in the measured velocity moments of low-mass galaxies (such as the projected velocity dispersion and the 4$^{\\rm th}$ order Gauss-Hermite coefficient $h_4$). It is shown that there is a noticeable change in the observed velocity dispersion if the intrinsic velocity dispersion of a galaxy is of the same order as the binary velocity dispersion. The kurtosis of the line-of-sight velocity distribution (LOSVD) is affected even at higher values of the intrinsic velocity dispersion.  Moreover, the LOSVD of the binary stars (i.e. the probability of finding a star in a binary system with a particular projected velocity) is given in closed form, approximating the constituent stars of all binaries to revolve on circular orbits around each other. With this binary LOSVD, we calculate the observed LOSVD, the latter quantifying the movement of stars along the line of sight caused both by the stars' orbits through the galaxy and by the motion of stars in binary systems. As suggested by the changes induced in the moments, the observed LOSVD becomes more peaked around zero velocity and develops broader high-velocity wings. These results are important in interpreting kinematics derived from integrated-light spectra of low-mass galaxies and many of the intermediate results are useful for comparison with Monte Carlo simulations. ",
        "introduction": "The usual approach to derive the kinematics (mean velocity along the line of sight, the velocity dispersion, \\ldots) of galaxies from observed spectra is comparing them to the spectra of so-called template stars. The broadening of the absorption lines in the galaxy spectra is then interpreted as a result of the orbital motion of the myriad of stars along the line of sight through the galaxy~: each star has a Doppler shifted spectrum and the absorption lines consequently appear at slightly different wavelengths. These template stars are carefully selected~: they should not rotate too rapidly, not be member of a binary system or be peculiar in any other way so as to be a good representation of the average stellar population. Thus, the broadening of the absorption lines is explained completely in terms of stellar orbital motions.  However, a large fraction of the stars in a galaxy are members of a binary (or even multiple) system. Stars in binary systems orbit the center of mass of the two stars and this extra velocity adds to the velocity dispersion of the galaxy. Hence, the measured kinematics of galaxies are not independent of the population of binary stars.  Previous authors have used Monte Carlo simulations to estimate the influence of binaries on the observed velocity dispersion (see Hargreaves {\\em et al.} \\cite{har} and references therein). This work focused mainly on observations of dwarf galaxies in the Local Group or globular clusters that can be resolved into individual stars. For a sample of stars in a galaxy, radial velocities are measured and from these velocities, the mean velocity and velocity dispersion profiles can be established. The Monte Carlo calculations by construction are very suited to interpret these ``discrete'' data, e.g. it is possible to check the effect of repeated observations to weed out short-period binaries. The crowded inner regions of globulars or dwarf galaxies outside of the Local Group however are impervious to this kind of study.  There, integrated-light spectra are needed to measure the kinematics.  As bigger telescopes make it possible to penetrate to ever lower levels of surface brightness and hence to study low-mass stellar systems at large distances, it is important to investigate the effect of binary stars on the derived LOSVD. Since Monte Carlo calculations at present are unable to yield the full LOSVD or even higher order moments such as the kurtosis without excessive computational effort, they cannot be employed in this context. Hence, we give analytical expressions for the velocity moments of the binary stars and for the LOSVD itself (approximating all binaries to consist of stars on circular orbits). ",
        "conclusions": "The maximum additional velocity dispersion due to binary stars is estimated at $\\sigma_{\\rm b} \\approx 3$~km/s, in good agreement with other authors. Only stellar systems with intrinsic velocity dispersions comparable to this value will have observed velocity dispersions that are noticeably affected by the presence of binaries, i.e. dwarf Spheroidals such as those found in the Local Group (Dekel~\\&~Silk \\cite{ds}), globular clusters, low-surface-brightness disk galaxies (Bottema \\cite{bot}) and the central regions of nucleated dwarf ellipticals. Not only the velocity dispersion but also the shape of the galaxy's LOSVD is altered. The presence of binaries has a measurable effect on the kurtosis of the observed LOSVD for stellar systems with a velocity distribution as high as 10~km/s.  For stellar systems with lower velocity dispersions, the LOSVDs will be even more distinctly non-Gaussian, being strongly peaked with broad wings. This feature can mimic radial anisotropy which combined with the enhanced velocity dispersion could lead to an over-estimation of the dark matter content of dSphs and globulars.  If the kinematics of a galaxy are derived from the radial velocities of discrete stars -- as is the case for the dwarf spheroidals in the Local Group -- repeated observations can eliminate the binaries from the star sample.  Methods that rely on integrated light spectra -- e.g. the crowded regions of globular clusters or dwarf galaxies outside the Local Group -- will be plagued by the above mentioned effects."
    },
    "0201/astro-ph0201297_arXiv.txt": {
        "abstract": "This article studies the structure of the Draco dwarf spheroidal galaxy with an emphasis on the question of whether the spatial distribution of its stars has been affected by the tidal interaction with the Milky Way, using R- and V-band CCD photometry for eleven fields.  The article reports coordinates for the center, a position angle of the major axis, and the ellipticity.  It also reports the results of searches for asymmetries in the structure of Draco. These results, and searches for a ``break'' in the radial profile and for the presence of principal sequences of Draco in a color-magnitude diagram for regions more than 50~arcmin from the center, yield no evidence that tidal forces from the Milky Way have affected the structure of Draco. ",
        "introduction": "\\label{intro}  This article studies the structure of the Draco dwarf spheroidal (dSph) galaxy.  The dSph galaxies are characterized by small size, low luminosity, low surface brightness, and old to intermediate age stellar populations.  The Local Group dSphs are clustered around and appear to be gravitationally bound to the much more luminous spiral galaxies (see the review by van den Bergh 2000). Draco, together with at least eight other dSphs, is a satellite galaxy of the Milky Way.  Draco is 80 $\\pm$ 7~kpc from the Sun (Aparicio, Carrera, \\& Mart\\'{i}nez-Delgado 2001, ACMD hereafter) and, although its tangential velocity is not known, the measured radial velocity, corrected for solar motion, of $-98$~km~s$^{-1}$ (Olszewski \\etal\\ 1995, Armandroff \\etal\\ 1995) implies that Draco is currently approaching the Milky Way.  Even though the masses of dSphs are small in comparison to those of luminous spiral and elliptical galaxies, many of their mass-to-light ratios ($\\cal{M}/\\cal{L}$\\,s hereafter) are very large (Aaronson 1983, Mateo 1998).  For example, the $\\cal{M}/\\cal{L}_{\\rm V}$ of Draco is 90 assuming that mass follows light (Armandroff \\etal\\ 1995) and 340 -- 610 using more realistic models (Kleyna \\etal\\ 2001) --- this is the highest measured value among the Galactic dSphs.  The presence of non-luminous, or dark, matter in Draco is the most direct interpretation of its large $\\cal{M}/\\cal{L}$.  However, several authors have proposed alternative explanations for the large measured $\\cal{M}/\\cal{L}$\\,s of dSphs in general and Draco in particular.  These explanations invoke either a modification of the laws of gravity (MOND, Milgrom 1983) or the dSph being far from virial equilibrium due to its interaction with the Galactic tidal field.  Kuhn \\& Miller (1989) and Kuhn (1993) proposed that a resonance between the orbital frequency and the frequency of internal collective oscillation modes of a dSph drives the dSph far from virial equilibrium.  In this picture, a dSph with a large measured $\\cal{M}/\\cal{L}$ is a gravitationally unbound stellar system which does not dissipate quickly because the stars are on Galactic orbits that keep them together.  However, Sellwood \\& Pryor (1998) show through numerical simulations that only a stellar system with even lower central concentration than those for the observed dSphs has collective modes that are not strongly damped and even in such a system these modes are not excited by coupling to the orbital motion.  Oh, Lin, \\& Aarseth (1995) modeled weak, but non-resonant, tidal interactions between the Galaxy and a dSph that led to the formation of tidal debris around the dSph but did not increase the velocity dispersion and $\\cal{M}/\\cal{L}$ to the values measured for real dSphs.  Piatek \\& Pryor (1995) examined whether a single strong tidal shock that disrupts a dSph can produce the large velocity dispersion and $\\cal{M}/\\cal{L}$ values measured for real dSphs.  This study found that such tidal interactions do not increase the velocity dispersion but instead they produce a large velocity gradient along the major axis.  Kroupa (1997) studied models in which the tidal debris from the dSph is aligned along the line of sight, in which case the velocity gradient masquerades as a large velocity dispersion. These models require that the dSphs with large measured $\\cal{M}/\\cal{L}$\\,s are on nearly radial orbits.  They also predict vertical broadening of the principal sequences, such as the horizontal branch, in the color-magnitude diagram (Klessen \\& Kroupa 1998, Klessen \\& Zhao 2001).  The presence of tidal debris around a dSph does not prove that their measured $\\cal{M}/\\cal{L}$\\,s have been raised by tidal interactions.  For example, Grillmair \\etal\\ (1995) and Leon \\etal\\ (2000) detected tidal debris around globular clusters, which have measured $\\cal{M}/\\cal{L}_{V}$\\,s of 1 -- 3 (Pryor \\& Meylan 1993) that agree well with those expected from their stellar populations. However, because the spatial extent and the surface density of tidal debris depends on the mass distribution, intrinsic $\\cal{M}/\\cal{L}$, and orbital elements of the dSph, along with the Galactic potential, detecting and quantifying the tidal debris can yield information about these quantities (\\textit{e.g.}, Kuhn 1993, Moore 1996, Johnston \\etal\\ 1999a, Johnston \\etal\\ 1999c).  In addition, deriving the amount and distribution of dark matter in a dSph using the kinematics of a tracer population requires measurements of the projected density profile of that population.  Limits on the central density of dark matter are particularly sensitive to the shape of the profile at large radii (Pryor 1994).  Irwin \\& Hatzidimitriou (1995; IH hereafter; see the references therein for earlier work) derived radial profiles and structural parameters for most of the Galactic dSphs using star counts from Palomar and UK Schmidt telescope plates.  IH determined limiting, or tidal, radii for a dSph by fitting single-component, isotropic King models to its profile.  They noted that in many cases the dSph has a profile that is above the fitted King model at large radii, which they interpret as evidence for ``extra-tidal'' stars.  Many subsequent studies have interpreted the IH tidal radius as the boundary between gravitationally bound and unbound populations.  However, the two-body relaxation time of every dSph is longer than its age (Webbink 1985) and so there is no reason that the dSph must resemble a King model. Indeed, it is possible to find many equilibrium models that fit perfectly almost any projected density and projected velocity dispersion profile (Dejonghe \\& Merritt 1992).  Several groups have recently studied the structure of the Draco dSph.  Smith, Kuhn, \\& Hawley (1997) detected apparent stars of Draco extending up to 3~degrees east of the center --- far beyond the 28.3~arcmin tidal radius determined by IH.  They interpret these stars as a tidally unbound population.  In contrast, Piatek \\etal (2001; hereafter P01) found evidence for Draco stars beyond the IH tidal boundary, though only weak evidence for stars with distances as large as 1~degree.  Odenkirchen \\etal (2001b; hereafter Od01) studied the morphology of Draco using Sloan Digital Sky Survey data for a wide region around the galaxy.  They derived a limiting (or tidal) radius for Draco of 49.5~arcmin, which is 75\\% greater than the IH value. Thus, they argue that the stars of Draco detected by P01 are within the tidal boundary and therefore are gravitationally bound to Draco. The Od01 upper limit for the surface density of Draco stars beyond their limiting radius is a factor of ten lower than the surface density detected by Smith, Kuhn, \\& Hawley (1997).  Od01 argue that this detection resulted from an incorrectly estimated background. Interestingly, they also note that an exponential model fits their data better than a King model, so the actual existence of a limiting radius is called into doubt.  Finally, ACMD derive a radial profile of Draco that is in broad agreement with that of Od01.  The large discrepancy in the values for the limiting radius of Draco obtained by IH and Od01 underscores the large uncertainties in this fitted parameter.  Since the dynamics of the tidal debris is decoupled from the internal dynamics of the dSph, the surface density profile of the tidal debris and that of the dSph should have different spatial structure.  Johnston \\etal\\ (1999b) performed numerical experiments which show that a robust indicator of tidal debris around a dSph is the presence of an abrupt change in the local power-law index of the radial surface density profile of a dSph -- a ``break'' in the profile.  The presence of tidal debris causes the surface density to decrease less steeply in the outer regions of a dSph.  Note, however, that Kroupa (1997) argues that a profile resembling that of a real dSph may be produced by an unbound population of stars.  It seems, therefore, that the limiting radius of a dSph is not a trustworthy representative of the tidal boundary of a dSph.  In this article we report the results of a study of the structure of Draco based on the R- and V-band photometric data for eleven fields in and around Draco.  We derive such structural parameters as the center, position angle, and ellipticity.  In addition, we search for tidal debris using methods based on color-magnitude diagrams, the radial profile of the surface density, and the shape of isopleths of a map of the surface density.  We compare our results to those from the existing studies of Draco.  Section~\\ref{data} describes the data and its reduction. Section~\\ref{cmds} presents the color-magnitude diagrams (CMDs) for the new N1 and S1 fields, describes the procedure of discriminating between non-members and possible members of Draco based on location in the CMD and image morphology, and compares the sample of possible members of Draco from this paper to those from ACMD and Od01. Section~\\ref{nomodel} derives model-independent structural parameters of Draco and a map of the projected density of the galaxy.  The section also comments on the reality of asymmetries in the projected density map.  Section~\\ref{discussion} summarizes and discusses our main results. ",
        "conclusions": "\\label{discussion}  The average center of Draco measured with three methods described in Sections~\\ref{xyp}\\ and \\ref{statexp}\\ is at $\\alpha$ = 17$^{\\rm h}$ 20$^{\\rm m}$ 18.1$^{\\rm s}$ and $\\delta$ = 57$^\\circ$ 55$^{\\prime}$ 13$^{\\prime\\prime}$ (J2000).  The uncertainty is about 0.1~arcmin in both coordinates.  This center is 40~arcsec east and 19~arcsec north of the center reported by Od01, a difference that is somewhat larger than that expected from the 25~arcsec uncertainty in right ascension and 11~arcsec uncertainty in declination of the Od01 value. However, the difference is smaller than twice the uncertainty and, thus, not statistically significant.  The position angle of the major axis of Draco is $90.6\\pm 4.6$~degrees based on the objects within approximately 10~arcmin of the center (see Section~\\ref{xyp}).   This value is in agreement with the $88\\pm 3$~degrees measured by Od01.  Fitting ellipses to the smoothed surface density, described in Section~\\ref{statexp}, gives a range of position angles consistent with the above values.  Od01 and this study find no evidence that the position angle of the major axis varies with semi-major axis.  The ellipticity of Draco, determined from the average of the values for the ellipses fitted to the smoothed surface density, is $0.331\\pm0.015$.  This average value is greater than the $0.29\\pm 0.02$ determined by Od01.  Our value is less reliable because of the problems with the photometry and completeness corrections in the N1 field described in Sections~\\ref{completeness} and \\ref{selecting}.  Tidal debris projected onto a bound dSph can produce a small asymmetry in the surface density map (Mayer \\etal\\ 2001).  A larger asymmetry might arise if the observed dSph consists primarily of unbound tidal debris (Kroupa 1997, Klessen \\& Kroupa 1998).  The contours of the smoothed surface density in Figure~\\ref{dcontour} show an apparent ``shoulder'' about 10~arcmin east of the center.  We tested the statistical significance of the apparent asymmetry along the major and minor axes and found that both can occur by chance 81\\% and 69\\% of the time, respectively.  Therefore, we find no compelling evidence for asymmetries in Draco, tidally induced or otherwise.  Figure~\\ref{allrp} shows that the radial profile of Draco from this study agrees with the radial profiles from IH and Od01 within the uncertainties.  The radial profile of Draco does not show evidence of an abrupt change, or break, in the slope. In addition, the cmd in Figure~\\ref{cmd-tb} does not show the principal sequences of Draco for the region beyond the Od01 tidal boundary, which is also about the last point in our radial profile.  Thus, we find no evidence that Draco is surrounded by tidal debris."
    },
    "0201/astro-ph0201542_arXiv.txt": {
        "abstract": "We present a variety of weak lensing results based on the ongoing analysis of $R_C$-band imaging data from the Red-Sequence Cluster Survey (RCS). We briefly discuss the weak lensing signal induced by intervening large scale structure (cosmic shear), and study the properties of the dark matter halos surrounding galaxies with $19.5<R_C<21$ (which have a median redshift of $z=0.35$) using a parametrized mass model for the galaxy mass distribution. This allows us for the first time to constrain the extent of the halos. We find a robust upper limit for the truncation parameter $s<470 h^{-1}$ kpc (99.7\\% confidence). We also study the biasing properties of these galaxies as a function of scale. We find that both the bias parameter $b$ and the galaxy-mass cross-correlation coefficient $r$ vary with scale (on scales $0.1-6~h^{-1}$ Mpc). Interestingly, we find that $r=0.57^{+0.08}_{-0.07}$ on scales $\\sim 0.5 h^{-1}$ Mpc. ",
        "introduction": "The Red-Sequence Cluster Survey\\footnote{\\tt http://www.astro.utoronto.ca/${\\tilde{\\ }\\!}$gladders/RCS} (e.g., Gladders \\& Yee 2000) is the largest area, moderately deep imaging survey ever undertaken on 4m class telescopes. The planned survey comprises 100 square degrees of imaging in 2 filters (22 widely separated patches imaged in $R_C$ and $z'$), and will provide a large sample of optically selected clusters of galaxies with redshifts $0.1<z<1.4$.  The survey allows a variety of studies, such as constraining cosmological parameters from the measurement of the evolution of the number density of galaxy clusters as a function of mass and redshift, and studies of the evolution of cluster galaxies, blue fraction, etc. at redshifts for which very limited data are available at present.  The data are also useful for a range of lensing studies. Strong lensing by clusters of galaxies allows a detailed study of their core mass distribution. In addition, given the shallowness of the survey, the arcs are sufficiently bright to be followed up spectroscopically (Gladders, Yee, \\& Ellingson 2001). Thanks to the large magnifications of the arcs, the survey allows unprecedented studies of the properties of high redshift galaxies. Furthermore, in combination with detailed modeling of the cluster mass distribution, the geometry of the images can be used to constrain $\\Omega_m$.  In these proceedings we concentrate on some of the weak lensing applications. A detailed description of the data and the weak lensing analysis is presented in Hoekstra et al. (2002a).  A careful examination of the residuals suggests that the object analysis, and the necessary corrections for observational distortions work well, which allows us to obtain accurate measurements of the weak lensing signal. ",
        "conclusions": ""
    },
    "0201/astro-ph0201068_arXiv.txt": {
        "abstract": "\\noindent We examined the fall-back disk models, and in general accretion, proposed to explain the properties of anomalous X-ray pulsars (AXPs), soft gamma repeaters (SGRs), and radio pulsars (PSRs). We checked the possibility of some gas remaining around the neutron star after the supernova explosion. We also compared AXPs and SGRs with the X-ray pulsars in X-ray binaries. We conclude the existing models of accretion from a fall-back disk are insufficient to explain the nature of AXPs/SGRs, particularly the SGR bursts. We also discussed the proposed model of combination of magnetic dipole radiation and propeller torques in order to explain the evolution of radio pulsars on the P-\\.{P} diagram. We found that the predictions of this model contradict the observational data. \\\\ \\\\ KEY WORDS AXP, SGR, PSR, SN ",
        "introduction": "In the last few years soft gamma ray repeaters (SGRs) and anomalous X-ray pulsars (AXPs), which are considered as a separate class of X-ray pulsars, have attracted much attention (see Mereghetti 2001 and references therein). The only important difference between SGRs and AXPs is that SGRs have active periods showing gamma ray bursts. A puzzling property of these objects is that their X-ray luminosities (L$_x$) are 1-3 orders of magnitude higher than their rate of rotational energy loss (\\.{E}). Besides, repeating gamma ray bursts of SGRs is the most important property of this class of objects. In order to explain the main physical peculiarities of these objects, not only a theory overstepping the limits of the standard NS/pulsar physics is needed, but also examining another type of NS, known as radio quiet NS, with low X-ray luminosity is required.  There is no radio radiation observed from AXPs/SGRs, only the upper limits for radio fluxes (and luminosities) are known (Gaensler et al. 2001). It does not necessarily mean that these objects do not radiate in the radio band only because of failing to observe radio pulses, because most of the pulsars are born with low luminosities (Lyne et al. 1998; Allahverdiev et al. 1997) and also because of the beaming factor. However, if somehow it is shown that AXPs and SGRs do not radiate in the radio band, then there will be a sharp drop in the radio luminosity of pulsars including AXPs and SGRs as we go to higher magnetic fields. It is difficult to check this possibility because it is not easy to increase the number of  such observational data as the number of AXPs and SGRs are small and they are located at large distances.  The so called magnetar model is proposed by Thompson and Duncan (1995) as a potential explanation for the nature of AXPs and SGRs. According to this model these sources are isolated rotating NSs with surface dipole magnetic fields of 10$^{14}$-10$^{15}$ G, extracting their luminosity from the decay of the magnetic field.  The alternative model is that the X-ray luminosity is due to accretion from a fall-back disk that is left over from the supernova that formed the NS (Chatterjee et al. 2000; Alpar 2001). In this work we will show the incompetence of the $existing$ accretion models in explaining the AXPs, SGRs, and also radio pulsar properties. In section 2, the similarities and differences between SGR bursts and flares of flare stars are shown qualitatively and quantitatively. In section 3, we investigated the possibility of a fossil disk remaining around the NS after the SN explosion, and if such a disk exists, the effects of this disk on the observed properties of the NS. In section 4, we compared the specific angular momentum transfer in AXPs and SGRs with the specific angular momentum transfer in X-ray binaries. In section 5, after a brief review of history of accretion theory, we showed that the proposed propeller mechanism for radio pulsars can not account for the evolution of radio pulsars on the P-\\.{P} diagram. The $existing$ fall-back disk models proposed to explain the nature of AXPs and SGRs are not well developed that some predictions of these models contradict some of the observational data. These will be shown below. ",
        "conclusions": "In this work, we have focused on if a residual (fossil) disk may remain around single NSs (AXPs, SGRs, and radio pulsars) after the supernova explosion and the effects of such a disk on the NS properties. We approached the problem from various different ways. We tried to clarify the physical nature of these objects and the existing accretion models in the light of observational data. In summary, we have showed that: \\\\ 1) The characteristics of SGR bursts are similar to the characteristics of flares of flare stars all of which have the same source (origin): the magnetic field of the star, but not accretion onto the star. \\\\ 2) After the supernova explosion, only a little amount of mass may remain around a single NS. The explosion energy of the thrown out matter is enough to sweep out most of the mass around the NS. \\\\ 3) The surrounding matter can also be swept out by angular momentum transfer due to magnetic braking and this is supported by observational data which do not show gaseous disks or planets (or the effects of these) near single pulsars with a magnetic field in excess of 10$^{10}$ Gauss. \\\\ 4) Even if such a little amount of mass remains around the NS, it spins up the NS (instead of spinning it down) unless the propeller mechanism is working together with accretion simultaneously. We showed this also by L$_x$/\\.{E} considerations in section 4. However, the propeller effect should be weak enough not to be observed since all the AXPs genetically connected to SNRs are pure shell type remnants. Fall-back disk models should be constructed taking these facts into account. \\\\ 5) Comparing the kinetic ages of pulsars with the age predictions of the propeller+dipole model for radio pulsars, we showed that there is a strong contradiction between them. So, propeller model does not work for radio pulsars.  $Acknowledgments$. It is a pleasure to thank Oktay H. Guseinov for fruitful and stimulating discussions. We thank Gennadi S. Bisnovatyi-Kogan for comments on the manuscript. We also thank T\\\"{U}B\\.{I}TAK, the Scientific and Technical Research Council of Turkey, for support through TBAG-\\c{C}G4."
    },
    "0201/astro-ph0201283_arXiv.txt": {
        "abstract": "We present a new grid (solar and halo abundance ratios) of state-of-the-art fully line-blanketed NLTE model atmospheres which covers the parameter range of central stars of planetary nebulae. The grid is available at the WWW. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0201/astro-ph0201410_arXiv.txt": {
        "abstract": "The radiative cooling  of electrons responsible for  the nonthermal synchrotron emission of large scale jets of radiogalaxies and quasars requires quasi continuous (in time and space) production  of relativistic electrons  throughout the jets over the scales exceeding 100 kpc.  While in  the standard paradigm  of large scale jets this implies {\\em in situ} acceleration of electrons,  in this letter we propose a principally different ``non-acceleration'' origin of these electrons, assuming  that they are implemented all over the length of the jet through effective development of electromagnetic cascades initiated by extremely high energy $\\gamma$-rays injected into the jet from the central object.  This  scenario provides a natural and very economic way to power the jets up to distances of 100 kpc and beyond. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0201/astro-ph0201189_arXiv.txt": {
        "abstract": "The cosmological recombination history of lithium, produced during Big--Bang nucleosynthesis, is presented using updated chemistry and cosmological parameters consistent with recent cosmic microwave background (CMB) measurements.  For the popular set of cosmological parameters, about a fifth of the lithium ions recombine into neutral atoms by a redshift $z\\sim 400$.  The neutral lithium atoms scatter resonantly the CMB at 6708~\\AA~ and distort its intensity and polarization anisotropies at observed wavelengths around $\\sim 300~\\mu$m, as originally suggested by Loeb (2001).  The modified anistropies resulting from the lithium recombination history are calculated for a variety of cosmological models and found to result primarily in a suppression of the power spectrum amplitude. Significant modification of the power spectrum occurs for models which assume a large primordial abundance of lithium. While detection of the lithium signal might prove difficult, if offers the possibility of inferring the lithium primordial abundance and is the only probe proposed to date of the large-scale structure of the Universe for $z\\sim 500-100$. ",
        "introduction": "Lithium is the heaviest stable element produced during Big--Bang nucleosynthesis \\citep*[and references therein]{bur01}. However, due to its low primordial abundance relative to hydrogen $X_{\\rm Li}\\sim 10^{-10}$--$10^{-9}$, its significance in the early Universe was thought to be restricted only to the formation of rare molecules such as LiH \\citep*[and references therein]{sta96}. Hence, the recombination history of lithium itself was only calculated as an intermediate step in a chain of chemical reactions and under a set of simplifying assumptions \\citep*[hereafter SLD96, SLD98]{pal95,sta96,sta98}. These preliminary calculations indicated that a substantial fraction ($\\ga 20$\\%) of the singly-charged lithium ions formed neutral atoms by recombination in the redshift interval $z\\sim 400$--$500$.  Recently, \\citet{loe01} has shown that the formation of neutral lithium can strongly modify the anisotropy maps of the cosmic microwave background (CMB) through the absorption and re-emission at its resonant 6708~\\AA~transition from the ground state.  Despite the exceedingly low lithium abundance\\footnote{Note that by the redshift of interest, all the $^7$Be produced during Big--Bang nucleosynthesis has been converted to $^7$Li.}  left over from the Big Bang, the resonant optical depth after lithium recombination is expected to be as high as $\\tau_{_{\\rm LiI}}\\sim 0.4 (X_{\\rm Li}/3.8\\times 10^{-10})$ at a redshift $z\\sim 400$ if half of the lithium ions recombine by then. The scattering refers to an observed wavelength of $\\lambda(z)=[6708~$\\AA$\\times (1+z)]=268.3~\\mu{\\rm m}\\times [(1+z)/400]$, where $X_{_{\\rm Li}}\\approx 3.8\\times 10^{-10}$ is the latest estimate of the lithium to hydrogen number density ratio \\citep{bur01}. \\citet{loe01} argued that resonant scattering would suppress the original anisotropies by a factor of $\\exp({-\\tau_{_{\\rm LiI}}})$, but will generate new anisotropies in the CMB temperature and polarization on sub--degree scales primarily through the Doppler effect.  Observations at different far--infrared wavelengths could then probe different thin slices of the early Universe. \\citet{zal02} calculated in detail the expected anisotropies in both the temperature and polarization of the CMB and assessed their detectability relative to the far-infrared background. They concluded that the modified polarization signal could be comparable to the expected polarization anisotropies of the far-infrared background on sub-degree angular scales ($\\ell \\ga 100$).  However, these calculations assumed a range of trial values for the neutral fraction of lithium in the redshift range $z=400$--500 lacking knowledge of its full redshift dependence.  In this paper, we calculate rigorously the recombination history of lithium and its subsequent effect on the CMB anistoropies. In \\S 2 we discuss the physics responsible for the recombination history of lithium and in \\S 3 we obtain the resulting optical depth for the 6708~\\AA~transition.  \\S 4 describes the effects of the lithium optical depth on the CMB power spectra. Finally, \\S 5 summarizes this work. ",
        "conclusions": "We have calculated the recombination history of lithium for a set of different cosmological parameters (see Table 1). The resulting opacity to the 6708~\\AA~ resonant transition of neutral lithium (Figs. 3 and 4) distorts both the temperature (Fig. 5) and polarization (Fig. 6) power-spectra of the CMB anisotropies over a band of observed photon wavelengths around $\\sim 300~\\mu$m.  The distortion is most pronounced in model E which assumes a relatively high lithium abundance of $X_{\\rm Li}=1.2 \\times 10^{-9}$. Our detailed results agree with earlier estimates by Loeb (2001) and Zaldarriaga \\& Loeb (2001).  The predicted polarization signal is comparable to the expected polarization anisotropies of the far-infrared background (FIB; see Zaldarriaga \\& Loeb 2001). A strategy for eliminating the contribution from the brightest FIB sources would be helpful in isolating the lithium signal.  The relevant wavelength range overlaps with the highest frequency channel of the {\\it Planck} mission ($352~\\mu$m), with the Balloon--borne Large--Aperture Sub--millimeter Telescope\\footnote{http://www.hep.upenn.edu/blast} (BLAST) that will have $250,\\ 350$ and $500~\\mu$m channels and with the proposed balloon--borne Explorer of Diffuse Galactic Emissions\\footnote{http://topweb.gsfc.nasa.gov} (EDGE) that will survey 1\\% of the sky in 10 wavelength bands between $230$--$2000~\\mu$m with a resolution ranging from $6^\\prime$ to $14^\\prime$ (see Table 1 in Knox et al. 2001).  However, in order to optimize the detection of the lithium signature on the CMB anisotropies, a new instrument design is required with multiple narrow bands ($\\Delta \\lambda/\\lambda\\la 0.1$) at various wavelengths in the range $\\lambda=250$--$350~\\mu$m. The experiment should cover a sufficiently large area of the sky so as to determine reliably the statistics of fluctuations on degree scales, and the detector should be sensitive to polarization. For reference, the experiment should also measure the anisotropies at shorter wavelengths where the FIB dominates. All of these requirements, of course, will make the detection of the lithium signature a difficult experiment. The difficulty cannot be fully accessed until more is known about the FIB and the sources responsible for it. Some considerations and strategies are discussed in \\citet{zal02}. Further, in order to detect the effect of lithium, high signal-to-noise maps of the primordial CMB at long wavelengths should be made for the same region of the sky.  Fortunately, these maps will become available from future CMB missions such as the {\\it Planck} satellite.  On the theoretical front, improvements on the predicted lithium recombination history can be made by performing an explicit calculation of the non-equilibrium populations of the excited states in the same level-by-level fashion as that of the hydrogen recombination calculation of \\citet{sea99,sea00}.  For hydrogen it was found that the excited state levels were overpopulated for $z < 800$ as the levels fall-out of equilibrium with the background radiation field. A slower photoionization rate coupled with faster cascade rates results in a faster net recombination rate ultimately resulting in a $\\sim$10\\% change in the ionization fraction. \\citet{sea99,sea00} also confirmed that all Lyman transitions are optically thick during hydrogen recombination indicating that Case B recombination holds at $z > 800$ but not at lower redshifts.  The optical depths computed here for lithium are less than one for the ground state transition suggesting that all other transitions should be optically thin, indicating Case A recombination. The Case A lithium equilibrium recombination rate coefficient of VF96 might result in an underestimation of the neutral lithium fraction (and the optical depth), but probably only by $\\sim 10\\%$.  We also note that non-conventional recombination histories may have a substantial impact on our predicted signal. For example, it has recently been proposed that recombination could have been delayed by the presence of additional radiation at $z\\sim 10^3$, possibly from stars, AGNs, or accretion by primordial compact objects \\citep*{pee00,mil01}. The neutral fraction of lithium (and its optical depth) could be larger than we calculated based on the standard recombination history, as the neutral lithium fraction is sensitive to the residual electron abundance. However, the stronger radiation field might override any gains in recombination of lithium due to the increase in its photoionization rate. Models with detailed radiation fields are needed to calculate the net effect that these processes have on the lithium optical depth.  Finally, while experiments to observe the lithium distortion on the temperature and polarization CMB power spectra appear to require redesign of detectors and new observational strategies, the benefits as mentioned in \\citet{zal02} could potentially be significiant: \\begin{enumerate} \\item If the lithium signal could be separated from the FIB contamination, it would offer the possibility of constraining the lithium primordial abundance. Inferring the lithium primordial abundance from stellar observations is known to be complicated by stellar lithium depletion and galactic lithium production. Further, the lithium primordial abundance is a sensitive indicator of the baryon abundance, its mean value as well as inhomogenities. It would then provide a possible means to discriminate among Big--Bang nucleosynthesis models. \\item The lithium signature on the CMB anisotropies is the only probe proposed so far for structure in the dark ages of the early Universe. Other methods, such as the suggestion by \\citet{ili02} that angular fluctuations in 21 cm emission from minihalos could probe redshifts between reionization and $z\\sim 20$, rely on the existence of collapsed objects and do not reach the high redshifts ($z\\sim$100 to $\\la 500$) probed by the lithium signal (see review by Barkana \\& Loeb 2001).  The lithium and 21 cm signals would both give measures of the baryonic density fluctuations, the latter during the era of early star formation leading to reionization and the former just before the condensation of the first objects. \\end{enumerate}"
    },
    "0201/astro-ph0201140_arXiv.txt": {
        "abstract": "A common proper motion survey of M dwarf stars within 8\\,pc of the Sun reveals no new stellar or brown dwarf companions at wide separations ($\\sim$\\,100-1400\\,AU). This survey tests whether the brown dwarf ``desert'' extends to large separations around M dwarf stars and further explores the census of the solar neighborhood.  The sample includes 66 stars north of $-30\\degr$ and within 8\\,pc of the Sun.  Existing first epoch images are compared to new $J$-band images of the same fields an average of 7 years later to reveal proper motion companions within a $\\sim$\\,4 arcminute radius of the primary star.  No new companions are detected to a $J$-band limiting magnitude of $\\sim$\\,16.5, corresponding to a companion mass of $\\sim$\\,40 Jupiter masses for an assumed age of 5\\,Gyr at the mean distance of the objects in the survey, 5.8\\,pc. ",
        "introduction": "Although the sub-stellar initial mass function (IMF) has been studied in a range of environments such as star-forming clusters (Luhman et al. 1998; Luhman 2000; Najita et al. 2000), young open clusters (Bouvier et al. 1998; Barrado y Navascues et al. 2001) and the field (Reid et al. 1999), the IMF of low-mass companions is not well understood, especially at ``wide'' ($>$\\,100\\,AU) separations. Radial velocity searches around solar-type main sequence stars (e.g., Mayor \\& Queloz 1995; Marcy \\& Butler 1996) have produced few confirmed brown dwarfs at separations $<$\\,3\\,AU.  Fewer than 0.5\\% of their sample have brown dwarf companions at those separations. A coronagraphic search for companions in the range 40-100\\,AU (Oppenheimer et al. 2001) produced only one brown dwarf, GJ 229B (Nakajima et al. 1995), well below the 17-30\\% multiplicity observed for all stars (Reid \\& Gizis 1997).  Other types of surveys, such as high spatial resolution space-based observations (Lowrance et al. 1999; Lowrance et al. 2000) and ground-based adaptive optics (Els et al. 2001), have also resulted in discoveries of low-mass stellar and sub-stellar companions.  However, the frequency of stellar and sub-stellar companions at close separations remains distinctly different, resulting in the idea that there is a ``brown dwarf desert''.  To date there has been only one systematic search for brown dwarf companions at wide separations and with a volume-limited sample (Simons et al. 1996; hereafter, SHK).  This was mainly a color-based search around M dwarfs within 8\\,pc of the Sun and did not turn up any new brown dwarfs, although, given the surprisingly blue colors of GJ 229B, cool brown dwarfs with intermediate $J$-$K$ colors may have been overlooked in the survey.  Proper motion searches for companions have been used for many years to identify low-mass objects (e.g., van Biesbroeck 1961) and offer a less biased way of finding low-mass companions than color-based surveys.  Therefore, we have conducted the planned second epoch survey of the SHK sample, in order to identify low-mass companions to M dwarfs at wide separations out to over 1000\\,AU. The choice of M dwarf primaries is significant:  Reid \\& Gizis (1997) and Reid et al. (1999) show that the distribution of mass ratios for a sample of 80\\% M dwarfs has a peak at $q=$ 0.95, where $q$ is the ratio of the secondary mass to the primary mass.  They conclude that their sample shows a distinct bias towards approximately equal-mass systems and that the mass function for stellar companions is different from the IMF of field stars.  If these conclusions extend to brown dwarf masses, M primaries may harbor more sub-stellar companions than other stellar types.  On the other hand, Reipurth \\& Clarke (2001) suggest that brown dwarfs have been ejected by dynamical interactions during the star formation process and cannot accrete enough mass to become stars. In this case M dwarf primaries may not be accompanied by such companions except in a multiple M dwarf systems with a correspondingly large gravitational potential.  Thus, our proper motion search around one spectral class of primaries fills a unique niche in the search for low-mass stellar and brown dwarf companions.  We describe the data acquisition and reduction in $\\S$\\,2 and discuss the results of the survey in $\\S$\\,3. ",
        "conclusions": ""
    },
    "0201/astro-ph0201376_arXiv.txt": {
        "abstract": "The annihilation rate of weakly interacting cold dark matter particles at the galactic center could be greatly enhanced by the growth of a density spike around the central supermassive black hole (SBH). Here we discuss the effects of hierarchical mergers on the central spike. Mergers between halos containing SBHs lead to the formation of SBH binaries which transfer energy to the dark matter particles, lowering their density. The predicted flux of annihiliation photons from the galactic center is several orders of magnitude smaller than in models that ignore the effects of SBHs and mergers.  Measurement of the annihilation radiation could in principle be used to constrain the merger history of the galaxy. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0201/astro-ph0201006_arXiv.txt": {
        "abstract": "I present  detailed models for  the formation of disk  galaxies, and investigate  which  observables  are  best  suited  as  virial  mass estimators.   Contrary to naive  expectations, the  luminosities and circular velocities  of disk galaxies are  extremely poor indicators of  total virial mass.   Instead, I  show that  the product  of disk scale length and rotation velocity squared yields a much more robust estimate.  Finally, I show  how this  estimator may  be used  to put limits  on  the efficiencies  of  cooling  and  feedback during  the process of galaxy formation. ",
        "introduction": "Currently, the  main uncertainties in our picture  of galaxy formation are related to the intricate processes of cooling, star formation, and feedback.   The  cooling  and  feedback  efficiencies  are  ultimately responsible  for setting  the  galaxy mass  fractions  $f_{\\rm gal}  = M_{\\rm  gal} / M_{\\rm  vir}$.  Here  $M_{\\rm gal}$  is the  total {\\it baryonic} mass of the galaxy (stars plus gas, excluding the hot gas in the halo) and $M_{\\rm vir}$ is the total virial mass.  Here I present new models for  the formation of disk galaxies, which I use to  investigate how well observables {\\it  extracted directly from these models} can be used to recover $f_{\\rm gal}(M_{\\rm vir})$.  Even though  the  assumptions  underlying  the model  are  not  necessarily correct, and  the phenomenological descriptions of  star formation and feedback are  certainly oversimplified, this  provides useful insights regarding the  ability of actual observations to  constrain the poorly understood astrophysical processes of galaxy formation. ",
        "conclusions": "We  have shown  that  the product  of  disk scale  length and  maximum rotation velocity squared  can be used as a  fairly accurate estimator of total virial mass. This in  turn can be used to obtain estimates of the galaxy mass  fractions as function of virial  mass, which contains important information about the  efficiencies of cooling and feedback. Another approach, that has been taken in the past, is to use published luminosity  functions and  luminosity-velocity relations  to construct halo velocity  functions (i.e., Newman  \\& Davis 2000;  Gonzales \\etal 2000; Kochanek  2001).  The main goal  of these studies  is similar to the work presented here, namely to circumvent the problems with poorly understood   astrophysical  processes   when   linking  the   observed properties  of galaxies  to those  of  their dark  matter halos.   Our results imply that great care is  to be taken in linking an observable velocity  such as $V_{\\rm  max}$ to  the circular  velocity of  a dark matter halo.  Based  on our results, we suggest  that the construction of a halo {\\it mass}  function using $M_{\\rm vir} \\propto R_d V^2_{\\rm max}$ may proof  more reliable. More details of  the results presented here can be found in van den Bosch (2002b)"
    },
    "0201/astro-ph0201230_arXiv.txt": {
        "abstract": "We present a {\\sl ROSAT} and {\\sl ASCA} study of the {\\sl Einstein} source X-9 and its relation to a shock-heated shell-like optical nebula in a tidal arm of the M81 group of interacting galaxies. Our {\\sl ASCA} observation of the source shows a flat and featureless X-ray spectrum well described by a multi-color disk blackbody model. The source most likely represents an optically thick accretion disk around an intermediate mass black hole ($M \\sim 10^2 {\\rm M}_{\\odot}$) in its high/soft state, similar to other variable ultraluminous X-ray sources observed in nearby disk galaxies. Using constraints derived from both the innermost stable orbit around a black hole and the Eddington luminosity, we find that the black hole is fast-rotating and that its mass is between $\\sim 20/({\\rm cos}\\ i)\\ {\\rm M}_{\\odot} - 110/({\\rm cos}\\ i)^{1/2} {\\rm M}_{\\odot}$, where $i$ is the inclination angle of the disk. The inferred bolometric luminosity of the accretion disk is $\\sim (8 \\times 10^{39} {\\rm~ergs~s^{-1}})/({\\rm cos}\\ i)^{1/2}$. Furthermore, we find that the optical nebula is very energetic and may contain large amounts of hot gas, accounting for a soft X-ray component as indicated by archival \\rosat\\ PSPC data. The nebula is apparently associated with X-9; the latter may be powering the former and/or they could be formed in the same event (e.g., a hypernova). Such a connection, if confirmed, could have strong implications for understanding both the birth of intermediate mass black holes and the formation of energetic interstellar structures. ",
        "introduction": "One of the most enigmatic X-ray-emitting objects is the source X-9 (1E 0953.8+6918), discovered with the {\\it Einstein Observatory} in the field close to the galaxy M81 (Fig.\\ 1; Fabbiano 1988). The source is located about 12\\farcm5 from the nucleus of the galaxy and 2$^\\prime$ from the galaxy's dwarf companion Ho IX ($1^\\prime$ corresponds to a projected separation of 1 kpc at the distance $D = 3.6$~Mpc; Freedman et al. 1994). The 0.2--4~keV flux of the source is only about a factor of $\\sim 2$ lower than the flux of the nucleus (a LINER). X-9 was also detected in subsequent \\rosat\\ and \\asca\\ observations at comparable flux levels. Clearly the source is not a transient. X-9 does, however, exhibit significant sporadic timing variability and therefore must primarily be a compact source (Immler \\& Wang 2001; Ezoe et al. 2001). Interestingly, Miller (1995) discovered that the source was projected inside a very  unusual \\ha-emitting nebula (Fig.\\ 1), which also emits strong [S{\\small II}] and [O{\\small I}] lines. He concluded that this nebula was shock-heated and might be a very energetic supernova remnant (SNR) or a superbubble. Furthermore, the nebula is apparently located within a nearly ``quiescent'' massive atomic and molecular gas concentration in a tidal arm (Concentration I, Yun et al. 1994; Fig.\\ 1a). The lack of a significant stellar population in this region has further led to the suggestion of the concentration being a protogalaxy (Henkel et al. 1993). The nature of this combination of stellar and interstellar features remains unknown.  \\begin{figure} \\centerline{\\hfil\\hfil \\psfig{figure=f1a.ps,height=3.4truein,angle=90.0,clip=} \\hfil\\hfil} \\centerline{\\hfil\\hfil \\psfig{figure=f1b.ps,height=3.2truein,angle=0.0 ,bbllx=152bp,bblly=395bp,bburx=474bp,bbury=715bp,clip=} \\hfil\\hfil} \\caption{\\protect\\footnotesize The field including M81 in near-UV ($\\sim 2490 $\\.A; Hill et al. 1992), with overlaid HI contours at 5, 10, 15, 20, and 25 $\\times 10^{20} {\\rm~cm^{-2}}$ (upper panel). A close-up of the region outlined by a box in the field of the protogalaxy is shown in the lower panel. The H$\\alpha$ image is provided by Miller (1995) and the \\rosat\\ HRI intensity contours are plotted at 0.24, 0.34, 0.44, and 0.54 $\\times  {\\rm~counts~s^{-1}~arcmin^{-2}}$. The structure of the X-ray emission is uncertain because of the complicated PSF of the observations at the large off-axis angle  (Immler \\& Wang 2001). } \\end{figure}  In this paper we present results from the analysis of a dedicated \\ASCA\\  and archival \\ROSAT\\ observations of the intriguing X-ray source X-9. We concentrate on the nature of the source and its relation to the environs, not on the detailed processes involved in the emission and evolution of the source. ",
        "conclusions": "The above spectral characteristics of X-9, together with its strong aperiodic variation (Immler \\& Wang 2001) and its lack of periodicity, argues against a scenario of the source as a Galactic compact object. Indeed, at the source's Galactic position of $l, b = 143^\\circ, 18^\\circ$, there is little room for such a scenario. Ezoe et al. (2001) have made additional arguments against the Galactic origin of X-9, based primarily on its high X-ray to optical emission ratio. In contrast, the model of X-9 as an accreting BH inside Concentration I at the M81 distance naturally accounts for the sporadic timing behavior, the convex-shaped spectrum, the excess absorption, and the large X-ray flux of the source.  But the apparent discrepancy between the GIS and PSPC spectral characteristics demands an explanation. The low $T_{\\rm in}$ as inferred from the PSPC spectrum is very abnormal, compared with similar ULXSs observed with \\asca. X-9 also shows strong flux variation: e.g., $2.5 \\times 10^{-12} {\\rm~ergs~s^{-1}~cm^{-2}}$ for the PSPC spectrum vs. $5.1 \\times 10^{-12} {\\rm~ergs~s^{-1}~cm^{-2}}$ for the GIS spectrum over the overlapping 0.5--2 keV band. But this variation is typical for ULXSs and is well within the range inferred from the archival GIS spectra of X-9. Indeed, the MCD model provides satisfactory fits to all individual GIS spectra, and $T_{\\rm in}$ is always greater than 1.2 keV. Clearly, the MCD model with the abnormal low $T_{\\rm in}$ cannot be a fair characterization of the GIS spectra. It is important to note that the MCD model has been tested primarily on X-ray spectra of heavily absorbed Galactic black hole candidates. It is possible that the model does not adequately describe the soft X-ray emission from the accretion disk of X-9.  Alternatively, the PSPC spectrum of X-9 may contain a significant soft X-ray contribution from diffuse hot gas. This may be expected in the region enclosed by the shock-heated optical nebula, or from relativistic particles accelerated in an expanding shock related to the nebula. The narrow bandwidth and very limited spectral resolution of the PSPC data, however, do not allow for a useful spectral decomposition of such a soft component from the uncertain disk emission of the source at the time of the observations. A spatial analysis of the archival \\rosat\\ PSPC and HRI data also shows evidence for the extended X-ray emission around the centroid of X-9 on scales comparable to the size of the nebula. But various systematic uncertainties (e.g., the source centroid shift caused by the off-axis instrumental point spread function) prevent us from a definitive measurement of the extended emission.  \\subsection{X-9 as an Accretting IMBH}  Assuming that the MCD model is correct, we can estimate the mass of the BH: $M_{\\rm BH} = (41 {\\rm M}_{\\odot}) \\alpha^{-1} (K_{\\rm MCD} /{\\rm cos}~i)^{1/2},$ where $\\alpha$ is the ratio of the inner disk radius to the last stable orbit radius of a non-rotating (Schwarzschild) BH, and both $K_{\\rm MCD}$ and $i$ are defined in the note to Table 1. Adopting $K_{\\rm MCD}$ in Table 1 gives $M_{\\rm BH} =  18/(\\alpha^2 {\\rm cos}~i)^{1/2} {\\rm M}_{\\odot}$. For a general rotating Kerr BH, $\\alpha \\gtrsim 1/6$ (e.g., Zhang et al. 1997) and thus $M_{\\rm BH} \\lesssim 108 {\\rm M}_{\\odot}/({\\rm cos}~i)^{1/2}$. The model, however, does not include various general relativistic effects (e.g., light bending). Ongoing modeling of such effects shows that the cosine law is approximately preserved for Schwarzschil BHs and that for extreme Kerr BHs the effective ${\\rm cos}\\ i = 0.17-0.4$ for $i = 5^\\circ-85^\\circ$ (Zhang et al. 2001). So the range of variation can be considerably smaller than the cosine law. Following Makishima et al. (2000), we can also estimate the bolometric luminosity of the disk as $L_{bol} = (8 \\times 10^{39} {\\rm~ergs~s^{-1}})/({\\rm cos}~i)^{1/2}.$ Because $L_{bol}$ should typically be smaller than, or at most comparable to, the Eddington luminosity  $1.3 \\times 10^{38} (M_{\\rm BH}/{\\rm M}_{\\odot})$, we obtain $\\alpha \\lesssim 0.76 ({\\rm cos}~i)^{1/2}$, indicating that the BH is fast-rotating ($\\alpha < 1$) and that $M_{\\rm BH} \\gtrsim 24 {\\rm M}_{\\odot}/{\\rm cos}\\ i$. The 90\\% statistical uncertainties in these mass limits are about 10\\%.  \\subsection{Optical Nebula}  The presence of the unusual shell-like optical nebula is an important part of the mystery about X-9 (Fig.\\ 1b). Optical spectroscopy of the nebula shows that its mean heliocentric velocity ($\\sim 47-52 {\\rm~km~s^{-1}}$) agrees with the velocity ($\\sim 45-65 {\\rm~km~s^{-1}}$) of the \\HI concentration (Adler \\& Westpfahl 1996; Miller 1995). The \\HI velocity field further shows that the concentration is part of the M81 group (Miller 1995). In contrast, the dwarf galaxy Ho IX has an optical heliocentric velocity of $119 \\pm 60$ (de Vaucouleurs et al. 1991) and is offset spatially from the centroid of the \\HI\\ concentration. Therefore, Ho IX may not be related to either the \\HI\\ concentration or the nebula (Fig.\\ 1a).  What might be the origin of the nebula around X-9? At the distance to M81, the nebula has an \\ha\\ luminosity $L_{{\\rm H}\\alpha} \\sim 1 \\times 10^{38} {\\rm~ergs~s^{-1}}$ (Miller \\& Hodge 1994), which is about three times more luminous than the bright SNR N49 in the LMC (Vancura et al. 1992). But the most distinct difference between the two is their sizes: $\\sim 250$ pc $\\times$ 475 pc for the X-9 nebula (Miller 1995) vs. 8 pc radius for N49. If the nebula is heated primarily by a shock (Miller 1995), a significant fraction (parameterized here as $\\xi$) of $L_{{\\rm H}\\alpha}$ may then be produced by the excitation as pre-shock hydrogen atoms drift into the post-shock region (e.g., Cox \\& Raymond 1985). Optical spectroscopy so far, however, places only an upper limit of $300 {\\rm~km~s^{-1}}$ on the shock velocity of the X-9 nebula  (Miller 1995). Assuming a shock velocity $v_2 \\gtrsim 2$ (in units of $10^2 {\\rm~km~s^{-1}}$ and the standard case B (that is, optically thick for all Lyman recombination lines), we can estimate from $L_{{\\rm H}\\alpha}$ the pre-shock neutral gas density as $n_o \\sim (10 {\\rm~cm^{-3}}) v^{-1}_2 R_2^{-2}\\xi$, where $R_2$ (in units of $10^2 {\\rm~pc}$) is the characteristic radius of the nebula. The kinetic energy is then $E_k \\sim (1 \\times 10^{53} {\\rm~ergs}) v_2 R_2\\xi$ if $n_o$ is spatially uniform. This is likely to be an overestimate because the bulk of the H$\\alpha$-emitting gas is expected to arise in dense filaments and clouds in which shock velocity tends to be low. Furthermore, a substantial fraction of the H$\\alpha$ emission may also be due to the ionization by the X-ray source (\\S 4.3). Nevertheless, comparisons with known SNRs suggest that the X-9 nebula appears to energetic to be due to a single normal supernova (SN).  Could the optical nebula then be a superbubble produced by multiple SNe and fast stellar winds of massive stars (e.g., Mac Low \\& McCray 1988)? At the position of X-9, there is indeed a fuzzy blue  object which may represent a stellar cluster ($m_{\\rm B} = 20$; Henkel et al. 1996; Miller 1995). If this is the case, the stellar cluster should then have an age $t_s \\gtrsim 10^7$ yr, as young stars do not appear to contribute much to the ionization of the optical nebula  and no far-UV emission peak appears at the position of X-9 (e.g., Fig.\\ 1). Comparing $t_s$ with the expansion age of the nebula $t_e \\sim 3R/5v = (1 \\times 10^6 {\\rm~yr}) R_2 v_2^{-1}$, we find  $v \\lesssim (5 {\\rm~km~s^{-1}}) R_2$. This is less than the turbulence velocity of the ISM and too small for shock-heating to effectively produce significant \\ha\\ emission. Therefore, unless being accelerated recently by an SN (or a hypernova; see later discussion) the optical nebula is probably not a superbubble created by a massive stellar cluster.  \\subsection{Association of the Optical Nebula and X-9}  Using the log(N)-log(S) function from the \\rosat\\ All Sky Survey (Hasinger et al. 1998), we estimate that the probability for a chance projection of an X-ray source comparable to, or brighter than, X-9 within a circle of $\\sim 10^{\\prime\\prime}$ radius is only about $\\sim 5 \\times 10^{-7}$. Therefore, X-9 is most likely associated with the nebula.  We speculate that the optical nebula may be directly related to the presence of X-9. One possibility is that they were born together. The nebula may be an interstellar remnant of a hypernova explosion that is substantially more energetic than a normal SN. Hypernovae have been postulated as the sources of some $\\gamma$-ray bursts observed at cosmological distances  (Paczy\\'nski 1998; Fryer \\& Woosley 1998). The proposed mechanism for hypernova explosions is the collapse of certain massive stars and/or their mergers with compact companions. Such an event leads to the formation of a BH and provides an extractable energy of $\\sim 10^{54}$~ergs (M\\'esz\\'aros, Rees, \\& Wijers 1999). Hypernovae may also be responsible for some \\HI supershells or holes observed in the interstellar medium (ISM) of nearby galaxies (Efremov, Elmegreen, \\& Hodge 1998; Loeb \\& Perna 1998), as well as relatively young energetic shell-like  nebulae (Wang 1999). The optical nebula around X-9 could well be such a hypernova remnant. The X-ray source may represent the resultant BH accreting from the materials falling back from the explosion. Although a binary may survive the explosion, the turning-on of a persistent accretion phase is expected to be long after the explosion remnant has disappeared. The soft X-ray contribution from the hypernova remnant may be significant in the PSPC spectrum, whereas the GIS spectrum is dominated by the bright accretion disk in a higher state, which makes the detection of the soft component difficult.  Another plausible scenario is that the nebula is currently powered by an intense outflow or a wind from X-9 as an accreting X-ray binary. If the nebula is purely caused by this outflow, the required mean energy output over the nebula's expansion age $t_e$ is then a few times $ 10^{39} {\\rm~ergs~s^{-1}}$, which is comparable to the X-ray luminosity of the source. In principle, the power of such an outflow could even be greater than the radiation luminosity of X-9, if it is similar to other types of accreting BH systems (i.e., Galactic micro-quasars or AGNs). In this case, the IMBH may be a Population III remnant (e.g., Madau \\& Rees 2001) and its companion may be formed from the collapse of surrounding molecular gas inside Concentration I (Henkel et al. 1993; Yun et al. 1994; Brouillet et al. 1992; Fig.\\ 1).  X-9 may also contribute to the ionization of the nebula. The MCD model of X-9 predicts an integrated ionizing photon rate of $\\sim 4 \\times 10^{40} {\\rm~s^{-1}}$ over the 0.016--0.2 keV range; photons at higher energies should mostly escape from the nebula. This rate is about a factor of 10 greater than that estimated for LMC X-1 (a stellar mass BH candidate), which is surrounded by an X-ray-ionizing nebula (Pakull \\& Angebault 1986). The nebula around X-9 shows indications of X-ray ionization. First, the H$\\beta$ to H$\\alpha$ flux ratio  ($\\sim 0.6$) of the nebula is high, indicating a high electron temperature of $\\sim 10^5$~K (Miller 1995).  Second, the shell-like morphology of the nebula appears to be rather diffuse, as is expected from the relative long absorption path-length of soft X-rays and its strong energy dependence (e.g., Rappaport et al. 1994).  Further scrutiny of the above scenarios and their relative importance is both desirable and possible. The nature of the blue object is yet to be determined: is it the optical counterpart of the accreting system or the outflow? The outflow may also be probed by observing its nonthermal radio emission, which should show a persistent flat or inverted radio spectrum (e.g., Fender 2001). The 20 cm radio continuum map of Bash \\& Kaufman (1986), which is centered on M81, does show a 3$\\sigma$ contour at the position of X-9. The total flux of $\\sim 1$ mJy, if indeed associated with the outflow, is substantially greater than that observed from Galactic micro-quasar-like objects. Detailed optical spectroscopy of the nebula will be especially important for determining the expanding velocity of the nebula and will provide important constraints on the extreme UV to soft X-ray radiation properties of X-9. Such radio/optical observations, combined with future high resolution X-ray imaging and spectroscopy, should allow for a firm determination of the relation between the nebula and X-9.  In summary, the association of X-9 with both the blue object and the shell-like $H\\alpha$ nebula as well as Concentration I provides an excellent opportunity for studying the nature of ultraluminous X-ray sources and for characterizing their radiation and outflow as well as their effects on the interstellar medium."
    },
    "0201/astro-ph0201094_arXiv.txt": {
        "abstract": "We study the effect of a prolonged epoch of reionization on the angular power spectrum of the Cosmic Microwave Background.  Typically reionization studies assume a sudden phase transition, with the intergalactic gas moving from a fully neutral to a fully ionized state at a fixed redshift. Such models are at odds, however, with detailed investigations of reionization, which favor a more extended transition.  We have modified the code {\\tt CMBFAST} to allow the treatment of more realistic reionization histories and applied it to data obtained from numerical simulations of reionization.  We show that the prompt reionization assumed by {\\tt CMBFAST} in its original form heavily contaminates any constraint derived on the reionization redshift.  We find, however, that prompt reionization models give a reasonable estimate of the epoch at which the mean cosmic ionization fraction was $\\approx 50 \\%$, and provide a very good measure of the overall Thomson optical depth.  The overall differences in the temperature (polarization) angular power spectra between prompt and extended models with equal optical depths are less than $1\\%$ ($10\\%$). ",
        "introduction": "At a redshift of $z\\approx 1100$ the intergalactic medium (IGM) recombined and remained neutral until the first sources of ionizing radiation formed.  While the distribution and evolution of these sources are unknown, the overall process of IGM reionization is fairly well understood, and can be divided into three phases.  First, individual HII regions developed around the sources. Then, these ionized regions grew in number and size until they overlapped, producing a sudden increase of the photon mean free path.  Finally, when all underdense regions and voids were completely ionized, photons penetrated into overdense clumps and filaments, bringing the reionization process to completion at a ``reionization redshift.''  It is clear that the mystery of the nature and evolution of the ionizing sources is intimately tied with the time scale over which these phases took place. Thus there is a great deal of physical information that could be gathered if the reionization redshift, $z_i$, and its duration, $\\Delta z$, can be firmly established.  Apart from the study of the (HI and HeII) Gunn-Peterson effect, which has not yet yielded conclusive results (Becker \\etal 2001; Gnedin 2001; Theuns \\etal 2001), measurements of Cosmic Microwave Background (CMB) anisotropies have the greatest potential for constraining these important quantities (Griffiths \\etal 1999). The scattering of CMB photons by free electrons damps the angular power spectrum of primary anisotropies by a factor of $e^{-2 \\tau}$ for large angular multipoles $\\ell \\gsim 100$ (Tegmark \\& Zaldarriaga 2000), where $\\tau$ is the Thomson optical depth. Using the currently available data, consistency with the lack of Gunn-Peterson trough and with the observed peak in the angular power spectrum at $\\ell \\approx 200$ (De Bernardis \\etal 2000; Hanany \\etal 2000; Padin \\etal 2001) is able to constrain $0.02 \\le \\tau \\le 0.44$ (De Bernardis \\etal 1997; Griffiths \\etal 1999; Tegmark \\etal 2001; Griffiths \\& Liddle 2001), although these results are somewhat dependent on the cosmological model assumed. Much better constraints are expected from polarization studies to be carried out by future satellites such as {\\it SPOrt}\\footnote{http://sport.tesre.bo.cnr.it} and {\\it PLANCK}\\footnote{http://astro.estec.esa.nl/SA-general/Projects/Planck}.  A second method of quantifying reionization is to examine $z_i$, and several recent studies have attempted to use the CMB to derive this redshift directly, using codes that assume a sudden epoch of reionization (e.g, Hu \\etal 95; Seljak \\& Zaldarriaga 1996).  For example, Schmalzing \\etal (2000) use {\\it MAXIMA} data in combination with cosmological parameters from independent measurements of Big Bang Nucleosynthesis and X-ray cluster data to constrain $z_i$. By performing a $\\chi^2$ analysis, they conclude that $z_i > 15 (8)$ at the 68\\% (95\\%) confidence level.  Similarly, Naselsky \\etal (2001), in addition to providing support to Schmalzing \\etal results, showed that polarization spectra are very sensitive to the reionization redshift.  An important caveat is implicit in these applications however, namely the assumption of prompt reionization.  While taking $\\Delta z = 0$ is a reasonable first step, more detailed simulations show that this approach paints a picture of reionization in only the broadest of strokes.  This is particularly true as the onset of reionization raises the overall temperature of the IGM, suppressing the further formation of objects that are too weakly bound gravitationally to overcome the drastic increase in thermal pressure (Barkana \\& Loeb 1999).  Thus recent numerical simulations of reionization (Ciardi \\etal 2000, hereafter CFGJ; Bruscoli \\etal 2000; Gnedin 2000; Benson \\etal 2000) have shown that the evolution of the mean ionization fraction is slow, has a nonlinear dependence on redshift, and occurs in an extremely patchy manner. All of these details may considerably affect the determination of $z_i$.  In this paper, we aim to clarify the effects of realistic reionization scenarios on the observables typically used to examine the reionization epoch.  To this end we have modified the most heavily relied on theoretical code for computing CMB fluctuations, {\\tt CMBFAST} (Seljak \\& Zaldarriaga 1996), to allow the treatment of such reionization histories, helping to clarify the best approach in trying to quantify and define the epoch of reionization.  The structure of this work is as follows.  In \\S 2 we discuss numerical models of reionization and how these have been incorporated into the {\\tt CMBFAST} code.  In \\S 3 we apply these techniques to study the impact of these models on the temperature and polarization spectrum of the CMB, and conclusions are given in \\S 4. ",
        "conclusions": "One of the first stages of nonlinear structure formation, reionization marked an important transition from a dark and relatively simple universe to one filled with a dazzling array of stars, galaxies, quasars, and other nonlinear objects.  And although one of our best probes of this transition is through the measurement of CMB fluctuations, the process of reionization itself is much more dependent on the complicated astrophysical issues important at low redshifts than the linear issues important at $z\\approx 1100.$  In this work, we have explored this transition, quantifying the impact of realistic scenarios of reionization on the angular power spectrum of the Comic Microwave Background.  While standard estimates assume prompt reionization, we have considered instead a range of simulated models, each with a prolonged reionization epoch.  We find that equating the redshift of full IGM ionization between these simulations and models that assume instantaneous reionization leads to widely discrepant temperature and polarization spectra. On the other hand, equating prompt and extended models with the same overall optical depth leads to differences in anisotropies that are nearly undetectable.  In this case the redshift of complete ionization is lost in the complicated details of the phase transition, and comparisons yield $z_i$ values corresponding to roughly the point of 50\\% ionization in the simulations, although even this value is model dependent.  It is clear then that while $z_i$ is useful as schematic tool, it is the total optical depth that is most accurate in providing a definition of the reionization epoch.      \\bigskip MB thanks M. Zaldarriaga for discussions and hospitality at IAS, Princeton where part of this work has been carried out. ES has been supported in part by an NSF MPS-DRF fellowship."
    },
    "0201/astro-ph0201431_arXiv.txt": {
        "abstract": "We report the discovery of two faint ultracool companions to the nearby (d$\\sim$17.9~pc) young G2V star HD~130948 (HR~5534, HIP~72567) using the Hokupa'a adaptive optics instrument mounted on the Gemini North 8-meter telescope. Both objects have the same common proper motion as the primary star as seen over a 7 month baseline and have near-IR photometric colors that are consistent with an early-L classification. Near-IR spectra taken with the NIRSPEC AO instrument on the Keck II telescope reveal K I lines, FeH, and H$_2$O bandheads. Based on these spectra, we determine that both objects have spectral type dL2 with an uncertainty of 2 spectral subclasses. The position of the new companions on the H-R diagram in comparison with theoretical models is consistent with the young age of the primary star ($<$0.8~Gyr) estimated on the basis of X-ray activity, lithium abundance and fast rotation. HD~130948B and C likely constitute a pair of young contracting brown dwarfs with an orbital period of about 10~years, and will yield dynamical masses for L dwarfs in the near future. ",
        "introduction": "The past decade of near-IR sky surveys and technological advances in high dynamic range imaging have found a large number ($\\sim$100) of very low-mass (VLM), ultracool objects. This has brought about spectral classification schemes (Burgasser et al. 2001; Geballe et al. 2001; Kirkpatrick et al. 1999, 2000; Mart\\'\\i n et al. 1997, 1998, 1999b) attempting to organize and understand them in the same way as we understand  main sequence stars through the MK spectral classification scheme. However, the interpretation of physical parameters from the classification schemes is a more complicated exercise with ultracool objects as the lack of a sustained hydrogen burning core creates a degeneracy between mass and age as the luminosity fades in time. Also, the spectra of these objects are significantly affected by dust in their atmospheres (Allard et al. 2001; Basri et al. 2000; Schweitzer et al. 2001) possibly introducing a weather-like time variable phenomenon (Bailer-Jones \\& Mundt 2001; Mart\\'\\i n et al. 2001; Nakajima et al. 2000). Just as stellar evolution theory was calibrated using the dynamical mass estimates of binary stars, it is important to check evolutionary tracks of VLM objects using low-luminosity binaries.  In recent years, there have been surveys using the high resolution capabilities of the HST (Mart\\'\\i n et al. 1999a, 2000a; Reid et al. 2001) and of large ground based telescopes (Close et al. 2002; Koerner et al. 1999; Mart\\'\\i n et al. 2000b) to look for companions to the known VLM objects. One goal of these searches is to build a sample of VLM binary systems in which accurate dynamical masses can be obtained. A handful of brown dwarf binaries are known, but only Gl~569B (Lane et al. 2001, Kenworthy et al. 2001), 2MASSW~J0746425+200032 (Reid et al. 2001), and 2MASSJ 1426316+155701 (Close et al. 2002) have periods $\\lesssim$10 years.  In this paper we add to the growing list of VLM ultracool binary systems. In a companion search around nearby, young (less than 1 Gyr), solar-type stars selected from the sample of Gaidos et al. (2000), we found two companions next to the star HD~130948. \\S2 overviews the observations, \\S3 presents the photometric, astrometric, and spectroscopic results which confirm that the companions are truly associated with the primary star.\\S4 discusses the placement of the objects on an HR-diagram compared with theoretical evolutionary models, and presents estimations of the age and mass of the companions. ",
        "conclusions": "HD~130948 is a chromospherically active single G2V star with high lithium abundance, and fast rotation (P=7.8 days). All these properties are indicative of youth (age$<$0.8~Gyr; Gaidos et al. 2000). The space motions of HD~130948 suggest that it could be related to the Ursa Major stream (Fuhrmann 2002, in preparation), which has an age of about 300~Myr.  The two new companions of HD~130948 are probably contracting brown dwarfs because of the young age of the primary star. With the aim of estimating their ages and masses, we placed HD~130948~B and C on an H-R diagram with theoretical evolutionary tracks shown in Figure~3. We used the evolutionary models of Chabrier et al. (2000) that include dust in the equation of state and the opacity because those are appropiate for L dwarfs (Allard et al. 2001). Basri et al. (2000) and Schweitzer et al. (2001) have estimated the effective temperatures (T$_{\\rm eff}$) of L dwarfs using the dusty models of Allard et al. (2001). Leggett et al. (2001) have used the same atmosphere models plus structural models for objects of known distance. We adopt T$_{\\rm eff}$=1950$\\pm$250~K for HD~130948B and C, which includes the whole range of T$_{\\rm eff}$ estimates for L0--L4 dwarfs in the literature. Our NIRSPEC data alone are not sufficient to tell whether HD~130948B and C have different L spectral type because the region that we observed does not contain features that are sensitive to changes in subclass in the range L0 to L4. We note, however, that if we force HD~130948B and C to lie on the same isochrone, their spectral types should differ by about 2 subclasses. Further spectroscopic observations, particularly at optical wavelengths, can test the agreement between the position of these objects in the H-R diagram and the model predictions.  For an age younger than 1~Gyr (consistent with youth of HD~130948A), the Chabrier et al. (2000) dusty models give a mass less than 0.075~M$_{\\odot}$ for HD~130948B, and less than 0.065~M$_{\\odot}$ for HD~130948C. It is very likely that both objects are young contracting brown dwarfs. For a total mass of the binary system of about 0.013~M$_{\\odot}$, and a semimajor axis of 2.4~AU, the orbital period should be $\\sim$10 years. Follow-up observations of this binary system over the next few years will yield dynamical masses for these two L dwarfs, which will extend the mass-luminosity-spectral type relation to cooler temperatures, and will provide two well constrained calibration points for the theoretical models describing low-mass, ultracool objects.  Although there are a handful of brown dwarfs known as companions to main sequence stars, HD130948 B-C is the first brown dwarf binary system imaged around a G-type star. This advance has been rendered possible by the high dynamic range provided by the Hokupa'a AO system on the Gemini-North telescope. At the time of writing this paper, 31 G-type stars less than 1Gyr old have been observed with Hokupa'a/Gemini in our ongoing survey for VLM companions to the stars in the Gaidos et al. (2000) sample and other nearby, young G-type stars. The survey observations are sensitive to objects 2 magnitudes fainter than the HD130948B-C objects at radii inbetween 10 AU and 100 AU from the stars. The detection of this new binary brown dwarf system in our survey gives a 3.2\\% $\\pm$3.2\\% frequency of brown dwarfs in the radius region we are sensitive to. This number is likely a lower limit because we are not sensitive to low mass brown dwarfs. Gizis et al. (2001) have reported a frequency of brown dwarf companions to G-type stars of 18\\% $\\pm$ 14\\% for separations larger than 1000 AU. Liu et al. (2002) have found an L-type companion at 14~AU of a G-type star using adaptive optics. Combining our result with that of Gizis et al. and Liu et al., we suggest that brown dwarf companions to G-type dwarfs with separations larger than 10~AU may be common. The brown dwarf desert may be restricted to separations less than 10~AU. This supports the theoretical models of Armitage \\& Bonnell (2002) that explain a lack of brown dwarfs within 10~AU of solar type stars as a consequence of orbital migration in circumstellar disks."
    },
    "0201/astro-ph0201161_arXiv.txt": {
        "abstract": "We report the first detection of pulsed X-ray emission from the young, energetic radio and $\\gamma$-ray pulsar \\psr. We find a periodic signal at a frequency of $f = 9.7588088 \\pm 0.0000026$~Hz (at epoch 51585.34104 MJD), consistent with the radio ephemeris, using data obtained withthe High Resolution Camera on-board the {\\it Chandra X-ray Observatory}.  The probability that this detection is a chance occurrence is $3.5 \\times 10^{-5}$ as judged by the Rayleigh test.  The folded light curve has a broad, single-peaked profile with a pulsed fraction of $23\\% \\pm 6\\%$. This result is consistent the \\ro\\ PSPC upper limit of $<18$\\% after allowing for the ability of {\\it Chandra } to resolve the pulsar from a surrounding synchrotron nebula. We also fitted {\\it Chandra} spectroscopic data on \\psr, which require at least two components, e.g., a blackbody of $T_{\\infty} = (1.66_{-0.15}^{+0.17}) \\times 10^6$~K and a power-law of $\\Gamma = 2.0 \\pm 0.5$.  The blackbody radius at the nominal 2.5 kpc distance is only $R_{\\infty} = 3.6 \\pm 0.9$~km, indicating either a hot region on a cooler surface, or the need for a realistic atmosphere model that would allow a lower temperature and larger area.  Because the power-law and blackbody spectra each contribute more than $23\\%$ of the observed flux, it is not possible to decide which component is responsible for the modulation in the spectrally unresolved light curve. ",
        "introduction": "\\label{sec:intro}  \\psr\\ is a 102 ms radio pulsar discovered by Johnston et al. (1992) and possibly associated with the supernova remnant \\snr\\ (McAdam, Osborne, \\& Parkinson 1993). With characteristic age $P/2\\dot P = 17,500$~yr and spin-down luminosity $\\dot{E} = 3.4 \\times 10^{36}$~ergs~s$^{-1}$ \\psr\\ is considered a young, energetic neutron star. The distance to the pulsar is somewhate uncertain.  According to the Taylor \\& Cordes (1993) free electron model of the Galaxy the pulsar lies 1.8~kpc away, however Koribalski et al. (1995) find a kinematic distance in the range $2.4-3.2$ kpc from H~I absorption.  Originally detected as the high-energy $\\gamma$-ray source 2CG342$-$02 by the {\\it COS B\\/} satellite (Swanenburg et al. 1981), \\psr\\ is one of approximately eight rotation-powered pulsars that are responsible for some of the brightest $\\gamma$-ray sources in the sky. Thompson et al. (1992; 1996) showed that photons detected from the direction of \\psr\\ by the EGRET instrument on board the {\\it Compton Gamma-Ray Observatory} are pulsed at the radio period.  All of the known $\\gamma$-ray pulsars are well observed at X-ray energies between 0.1 and 10~keV, where both surface thermal emission and magnetospheric synchrotron emission can be studied.  However, {\\it pulsed\\/} X-ray emission has not been detected from \\psr\\ despite several searches. This failure has hindered attempts to understand its X-ray emission mechanism(s).  Using an observation of \\psr\\ by the \\ro\\ PSPC, Becker, Brazier, Trumper (1995) placed an upper limit of 18\\% on the pulsed fraction.  They attributed this negative result to the diluting effect of an as-yet unresolved synchrotron nebula, recalling the history of Vela (\\\"Ogelman, Finley, \\& Zimmerman 1993), a pulsar of similar age and spin-down luminosity. Finley et al. (1998), using data from the \\asca\\ GIS, placed an upper limit of 22\\% on the pulsed fraction in the 2$-$10~keV band, although their image was severely contaminated by stray light from the bright low-mass X-ray binary 4U1705$-$44 just outside the field of view.  Finley et al. also analyzed a \\ro\\ HRI observation, from which they concluded that $57 \\pm 12$\\% of the photons associated with \\psr\\ actually come from a compact nebula with an exponential scale length of $\\approx 27^{\\prime\\prime}$.  When folding only the photons from the HRI point source at the radio period, they derived an upper limit of 29\\% on its pulsed fraction. An observation by the non-imaging {\\it Rossi X-ray Timing Explorer} performed during a low state of 4U1705$-$44 also failed to detect pulsations in the $9-18.5$~keV band. (Ray, Harding, \\& Strickman 1999). Here we report on {\\it Chandra} observations of \\psr\\ in which we detect its pulsed emission in X-rays. ",
        "conclusions": "\\label{sec:disc}  There are at least two mechanisms for the emission of broad-band pulsed X-rays from young rotation-powered pulsars like \\psr. One is non-thermal magnetospheric synchrotron from relativistic electrons and positrons created either in regions above the neutron star polar caps or in outer gaps.  The second is thermal emission from the hot surface, a result of initial cooling of the hot neutron star or reheating of the polar caps by back-flowing accelerated particles. Sometimes the shape of the pulse sheds additional light on the X-ray emission mechanism.  Sharp, narrow pulses of high amplitude can only be produced by a highly beamed, thus relativistic population of electrons, while quasi-sinusoidal pulses of low amplitude such as describe \\psr\\ can be produced by either mechanism. For those intermediate-age pulsars ($10^4 - 10^6$~yr) that are also EGRET sources, the presence of both types of X-ray source is usually discovered when spectrally resolved timing data are available (e.g., Wang et al. 1998; Pavlov et al. 2001). Unfortunately, in this case, the {\\it Chandra} HRC has little or no energy resolution.  Since each spectral component fitted to the ACIS spectrum contributes more than 23\\% of the flux in the total energy band to which the HRC is sensitive (60\\% from blackbody, 40\\% from power law), the pulsed fraction alone does not reveal the source of the pulsed X-rays.  Either or both components may contribute to the modulation. Future spectrally resolved X-ray observations with high throughput and moderate time resolution, such as with {\\it XMM-Newton}, could resolving this ambiguity.  The $>100$~MeV luminosity of most $\\gamma$-ray pulsars is a significant fraction of their spin-down power.  In the case of \\psr\\, this fraction is $\\approx 0.20$ if isotropic, while its $0.5-8$~keV non-thermal X-ray luminosity is only $1.3 \\times 10^{-4}\\,\\dot E$ including both pulsar and nebula, similar to that of other pulsars. Thompson et al. (1996) parameterized the EGRET spectrum of \\psr\\ as a broken power law, with photon index $\\Gamma = 1.27 \\pm 0.09$ from 50 MeV to 1~GeV, steepening to $\\Gamma = 2.25 \\pm 0.13$ above 1~GeV. Since there is no evidence for any {\\it unpulsed} $\\gamma$-ray emission, the EGRET spectrum can be compared directly with the power-law component of the pulsar point source in the {\\it Chandra} ACIS spectrum.  When extrapolated back to 10 keV, the EGRET spectrum nearly matches the X-ray flux, although the power-law X-ray slope itself, with $\\Gamma = 2.0 \\pm 0.5$, is somewhat steeper than the EGRET value. The non-thermal X-rays are not likely to be emitted by the same population of electrons/positrons as produce the $\\gamma$-rays, but may instead originate from a much less energetic population, such as secondary pairs from the inward emitted $\\gamma$-rays converting on the strong $B$-field near the neutron star surface.  This distinction is also supported by the observed phase offset between the X-ray and $\\gamma$-ray pulses (\\S 3).  The Wang et al. (1998) model predicts $L_x$ (2--10~keV) $\\approx 2 \\times 10^{31}$ ergs~s$^{-1}$ and $\\Gamma = 1.5$ from this process for \\psr, which is similar to its observed non-thermal X-ray component.  If its {\\it thermal\\/} luminosity comes partly from a small surface area, then inward flowing primary electrons impacting the polar caps may be responsible.  Wang et al. predicted $L_x \\approx 1 \\times 10^{33}$ ergs~s$^{-1}$ of thermal emission from this process, not far from the observered thermal luminosity of \\psr."
    },
    "0201/astro-ph0201482_arXiv.txt": {
        "abstract": "{Deep optical CCD images of the supernova remnant \\gsnr~were obtained and faint emission has been discovered. The images, taken in the emission lines of \\hnii, \\sulfur~and \\oiii, reveal filamentary structures in the east, south--east area, while diffuse emission in the south and central regions of the remnant is also present. The radio emission in the same area is found to be well correlated with the brightest optical filament. The flux calibrated images suggest that the optical filamentary emission originates from shock-heated gas (\\sulfur/\\ha\\ $>$ 0.4), while the diffuse emission seems to originate from an \\HII\\ region (\\sulfur/\\ha\\ $<$ 0.3). Furthermore, deep long--slit spectra were taken at the bright \\oiii\\ filament and clearly show that the emission originates from shock heated gas. The \\oiii\\ flux suggests shock velocities into the interstellar $``$clouds'' greater than 100 \\vel, while the \\siirat\\ ratio indicates electron densities $\\sim$240 cm$^{-3}$. Finally, the \\ha\\ emission has been measured to be between 7 to 20 $\\times$ \\flux. ",
        "introduction": "The Galactic supernova remnants (SNRs) have been identified by both radio (non-thermal synchrotron emission) and optical (optical emission lines) surveys. New searches in both wavebands continue to identify galactic SNRs (Fesen \\& Hurford 1995, Fesen et al. 1997; Green 2001 and references therein; Mavromatakis et al. 2000, 2001, 2002) but since the last few years, observations in X-rays have also detected new SNRs (e.g. Seward et al. 1995). The ratio of \\sulfur/\\ha~has become the standard discriminator used in optical SNR observations because the photoionized nebulae (like \\HII~regions and planetary nebulae) usually exhibit ratios of about 0.1-0.3, while collisionally ionized nebulae (like known Galactic SNRs) show ratios typically greater than 0.4 (Smith et al. 1993). Fesen et al. (1985) suggested that a division at \\sulfur/\\ha~$<$ 0.5 does not provide clear evidence to distinguish SNRs from photoionized regions and additional observations of the strong forbidden oxygen lines (\\oi, \\oii~and \\oiii) are needed to give a complete diagnostic. Furthermore, theoretical shock models , generally predicted \\sulfur/\\ha~ratios of 0.5 to 1.0 for SNRs (Raymond 1979, Shull \\& McKee 1979). \\par G 17.4$-$2.3 is not a well known SNR, and was first detected by Reich et al. (1988) in their Effelsberg 2.7--GHz survey, while its radio image was published by Reich et al. (1990). It is classified as a circular supernova remnant with an incomplete radio shell, characterized by diffuse shell--like emission, an angular size of $\\sim$24\\arcmin\\ and a radio spectral index of $\\sim$0.8 (Green 2001). Case \\& Bhattacharya (1998) calculated its surface brightness to be $1.3 \\times 10^{-21}$ W m$^{-2}$ Hz$^{-1}$ m$^{-1}$. Because, there is no direct distance determination, they have made an estimation by utilizing the radio surface brightness -- diameter relationship ($\\Sigma-D$) and found a distance of 8.5 kpc, but still the uncertainties are large ($\\sim$40\\%). Green et al. (1997), through their survey with the Parkes 64 m telescope, detected maser OH (1720 MHz) emission.  In radio surveys of the surrounding region, no pulsar was found to be associated with \\gsnr~but another SNR has been discovered in its neighborhood. G 17.8$-$2.6 has a very well defined shell, it lies about 30\\arcmin~north--east of \\gsnr~and has an angular diameter of 24\\arcmin~(Reich et al. 1988). Neither of these remnants has been detected optically in the past and from our observations no optical emission has been found in G 17.8--2.6. On the other hand, X--ray emission was not detected from \\gsnr~in the ROSAT All--sky survey, while there is some evidence of X--ray emission from the neighboring SNR G 17.8--2.6.  In this paper, we report the discovery of faint optical filaments from \\gsnr. We present \\hnii, \\sii~and \\oiii~images which show filamentary structure along the south--east edge of the remnant correlated very well with the radio emission. Spectrophotometric observations of the brightest filament were also obtained and the emission lines were measured. In Sect.2, we present informations concerning the observations and data reduction, while the results of our imaging and spectral observations are given in Sect. 3 and 4, respectively. In the last section (Sect. 5) we discuss the physical properties of \\gsnr. ",
        "conclusions": "The faint supernova remnant \\gsnr\\ was observed for the first time in major optical emission lines. The images show filamentary and diffuse emission structures. The bright \\oiii\\ filament is very well correlated with the remnant's radio emission at 1400 and 4850 MHz suggesting their association. The flux calibrated images and the long--slit spectra indicate that the emission arises from shock heated gas. The observed optical filamentary structure provides some evidence for significant inhomogeneities in the ambient medium, implying that the main blast wave propagates into an inhomogeneous medium."
    },
    "0201/astro-ph0201211_arXiv.txt": {
        "abstract": "{ We present an analysis of low- and medium resolution spectra of the very fast nova, Nova Cygni 2001/2 (V2275~Cyg) obtained at nine epochs in August, September and October, 2001. The expansion velocity from hydrogen Balmer lines is found to be 2100 km~s$^{-1}$, although early H$\\alpha$ profile showed a weak feature at $-$3500 km~s$^{-1}$, too. The overall appearance of the optical spectrum is dominated by broad lines of H, He and N, therefore, the star belongs to the ``He/N'' subclass of novae defined by Williams (1992). Interstellar lines and bands, as well as $BV$ photometry taken from the literature yielded to a fairly high reddening of $E(B-V)=1\\fm0\\pm0\\fm1$. The visual light curve was used to deduce M$_{\\rm V}$ by the maximum magnitude versus rate of decline relationship. The resulting parameters are: $t_0=2452141.4^{+0.1}_{-0.5}$, $t_2=2.9\\pm0.5$ days, $t_3=7\\pm1$ days, M$_{\\rm V}=-9\\fm7\\pm0\\fm7$. Adopting these parameters, the star lies between 3 kpc and 8 kpc from the Sun. ",
        "introduction": "Nova Cygni 2001/2 was discovered by A. Tago on two T-Max 400 films taken on Aug. 18 at 8\\fm8 apparent brightness. One day before discovery, nothing was visible at the nova location down to 12 mag (Nakamura et al. 2001). The spectroscopic confirmation was given by subsequent optical spectroscopy revealing hydrogen Balmer emission lines with deep P Cygni profiles. The H$\\alpha$ line profile suggested an early expansion velocity of 1700 km~s$^{-1}$ (Ayani 2001). The nova was also discovered independently by K. Hatayama (Nakano et al. 2001). A possible progenitor with USNO red magnitude 18\\fm8 and blue mag magnitude 19\\fm6 was identified by P. Schmeer (Schmeer et al. 2001).  Early photometric data consist of photographic and CCD photometric data published in the IAU Circulars. The full light curve can be reconstructed using visual data collected by the VSNET group\\footnote{\\tt http://www.kusastro.kyoto-u.ac.jp/vsnet}. The visual maximum occured at 6\\fm8 on 2001 Aug. 19.9 UT. Simultaneous color measurements were published by Sostero \\& Leopardo (in Nakano et al. 2001), who gave $B-V$=1\\fm1 in the maximum, suggesting substantial reddening. The color remained around $B-V\\sim 1\\fm0$  during the first week after the maximum, when the apparent brightness decreased down to 9\\fm2 (Samus et al. 2001).  The main aim of this paper is to present optical and far red spectra taken after the maximum, between $\\Delta$t = +2.3 d to $\\Delta$t = +59.1 d. The low- and medium resolution spectra were used to determine the main outburst properties, the expansion velocity and the interstellar reddening. We also make use of all publicly available visual photometric data collected by the VSNET group to estime the rates of decline and to check the photometric phases of the obtained spectra. ",
        "conclusions": "Based on the presented spectroscopic behavior of V2275~Cyg we identify the star as a member of the ``He/N'' subclass of classical novae following the definition by Williams (1992). All of the defining characteristics are satisfied: broad lines (HWZI$>$ 2500 km~s$^{-1}$) dominate the spectrum, the prominent lines are flat-topped with little absorption, a few forbidden lines occurred and $F({\\rm He II 4686})\\geq F({\\rm H\\beta})$. On the other hand, the very fast decline is also typical among the ``He/N'' novae. According to Williams (1992), this means that the broader lines of the ``He/N'' emission spectrum originate in a discrete shell ejected at high velocities from the white dwarf surface at the peak of the outburst. The stable system of narrow components in the H$\\alpha$ and O I profiles support the presence of one or more discrete shells, as provided by the emission line profile calculations of Gill \\& O'Brian (1999). An alternative to the rings+polar caps geometry is inhomogeneous ejecta with knots or clumps (see, e.g., Shore et al. 1997).  Since typical recurrent novae are classified as ``He/N'' novae, it is a natural consequence that the overall spectral appearance of V2275~Cyg is very similar to some well-observed recurrent novae. The most recent example is CI~Aql (Kiss et al. 2001), where the gross spectral characteristics are very similar. Furthermore, close similarity is evident from a comparison with other recurrent novae, such as U~Sco (Munari et al. 1999) or V394~CrA (Sekiguchi et al. 1989). The question arises: could V2275~Cyg be a recurrent nova, observed in outburst for the first time? We consider this to be quite unlikely. If the progenitor is indeed the star identified by Schmeer et al. (2001) with 18\\fm8 USNO-A2.0 red magnitude, and our calculated absolute magnitude is approximately correct, then the resulting absolute magnitude for the progenitor is M$_{\\rm vis}\\sim2-3$ mag. This excludes the possibility of a red giant secondary as usual in recurrent novae. To fit the suggested progenitor with a typical red giant star, the outburst should have a visual absolute magnitude over $-$12 mag, much brighter than any nova outburst ever observed. If the progenitor was fainter and hidden by the suggested USNO-A2.0 star, then the recurrent nova status is even less likely.  As expected from the very fast decline, V2275~Cyg closely resembles in some aspects the well-studied fast nova, V1500~Cyg. Strittmatter et al. (1977) presented a comparison of H$\\alpha$ and O I 8446 profiles, where they found a similar agreement between the line profiles to those presented in Fig.\\ 4. They concluded that the O I 8446 line is due to Lyman $\\beta$ fluorescence in clouds with high H$\\alpha$ optical depth. The strong resemblence suggest similar explanation in the case of V2275~Cyg.  Further photometric as well as spectroscopic observations are expected to extend the data baseline yielding to a better understanding of V2275~Cyg. At present, neither the unambiguously identified progenitor, nor the orbital period of the binary is known -- both are crucial for a reliable description of the system. On the other hand, theoretical light curve modelling may place stronger constraints on the luminosity of the outburst."
    },
    "0201/astro-ph0201027_arXiv.txt": {
        "abstract": "{ We present the first calibrated CCD images of the faint supernova remnant \\gsnr\\ in the emission lines of \\oii, \\oiii, \\hnii\\ and \\sii. The deep low ionization CCD images reveal diffuse emission in the south and central areas of the remnant. These are correlated with areas of intense radio emission, while estimates of the \\sii/\\ha\\ ratio suggest that the detected emission originates from shock heated gas. In the medium ionization image of \\oiii\\ we discovered a thin filament in the south matching very well the outer radio contours. This filament is not continuous over its total extent but shows variations in the intensity, mainly in the south--west suggesting inhomogeneous interstellar clouds. Deep long--slit spectra were also taken along the \\oiii\\ filament clearly identifying the observed emission as emission from shock heated gas. The \\ha\\ emission is a few times \\flux, while the variations seen in the \\oiii\\ flux suggest shock velocities into the interstellar clouds around or below 100 \\vel. The sulfur line ratio approaches the low density limit implying electron densities less than $\\sim$500 \\dens. ",
        "introduction": "The vast majority of supernova remnants have been discovered by their non--thermal synchrotron radio emission. Since the optical wavelengths suffer significantly more attenuation than the radio, the detection of a \\snr\\ in the optical band is a difficult task. The \\snr\\ \\object{G 114.3+0.3} was initially detected in the 21 cm continuum survey of Kallas \\& Reich (\\cite{kal80}) and subsequently studied in more detail by Reich \\& Braunsfurth (\\cite{rei81}). The remnant shows up in the radio as an ellipsoidal shell occupying an angular extent of $\\sim$60\\arcmin\\ $\\times$ 78\\arcmin. The radio emission is stronger in the south--east and south--west compared to other areas of the remnant. The surface brightness is low, while at the same time the source displays a high degree of polarization. More interest was  raised for this system following the proposal of F\\\"urst \\et\\ (\\cite{fur93}) and Kulkarni \\et\\ (\\cite{kul93}) that the pulsar PSR 2334$+$61 is associated with \\object{SNR 114.3+00.3}. This pulsar rotates at a period of 495 ms and its spindown rate suggests an age of $\\sim$ 41000 yr. The distance estimates to this pulsar--remnant association lie in the range of 2--3 kpc. ROSAT PSPC pointed observations showed that the pulsar is a weak X-ray source, while no emission was detected from the remnant itself (Becker \\et\\ \\cite{bec96}). \\par Fesen \\et\\ (\\cite{fes97}) reported the detection of faint filamentary structures in \\ha\\ in the west, south--west areas of the remnant. However, no images were shown due to the faintness of the filaments. Even though the spatial correlation of the optical emission with the radio contours favors their association, flux calibrated images in \\ha\\ and \\sii\\ or long--slit spectra are needed to establish the nature of the optical flux. In an effort to deepen our knowledge of the properties of the optically detected remnants, especially the faintest ones, we performed deep CCD imaging and spectral observations of \\gsnr. Information about the observations and the data reduction is given in Sect. 2. In Sect. 3 and 4 we present the results of our imaging observations and the results from the long--slit spectra taken at specific locations of interest. Finally, in Sect. 5 we discuss the physical properties of the remnant. ",
        "conclusions": "The supernova remnant \\gsnr\\ was observed in major optical lines. New diffuse structures were detected in the south--east and central areas of the remnant. The flux calibrated images imply that the diffuse emission is associated with \\gsnr. In addition, we discovered an \\oiii\\ filament in the south which is very well correlated with the radio contours at 4850 MHz. Long--slit spectra along this filament suggest that the emission arises from shock heated gas, while the sulfur line ratios indicate low electron densities."
    },
    "0201/astro-ph0201033_arXiv.txt": {
        "abstract": "I have used the Hipparcos Input Catalog, together with Kurucz model stellar atmospheres, and information on the strength of the interstellar extinction, to create a model of the expected intensity and spectral distribution of the local interstellar ultraviolet radiation field, under various assumptions concerning the albedo $a$ of the interstellar grains.  (This ultraviolet radiation field is of particular interest because of the fact that ultraviolet radiation is capable of profoundly affecting the chemistry of the interstellar medium.)  By comparing my models with the observations, I am able to conclude that the albedo $a$ of the interstellar grains in the far ultraviolet is very low, perhaps $a=0.1$.  I also advance arguments that my present determination of this albedo is {\\em much more reliable} than any of the many previous (and conflicting) ultraviolet interstellar grain albedo determinations. Beyond this,  I show that the ultraviolet background radiation that is observed at high galactic latitudes {\\em must be extragalactic} in origin, as it cannot be backscatter of the interstellar radiation field. ",
        "introduction": "The physics of the local interstellar medium is profoundly influenced by the intensity and spectrum of the local interstellar radiation field, particularly in the ionizing ultraviolet; but our knowledge of the {\\em strength} of that radiation field has considerable uncertainty.  There are few observations, and existing models are far from perfect.  In the present paper I improve on the method for estimating the brightness of the radiation field that was invented by Henry (1977), who showed that by simply integrating the flux that is expected from the stars that are included in the Yale Bright Star Catalog, one could provide a useful estimate of the local interstellar ultraviolet radiation field intensity.  Henry showed that, in the ultraviolet, such an integration {\\em converged}, demonstrating that the Bright Star Catalog went sufficiently deep to provide a complete result.  Henry's approach has the serious defect that, as he implemented it, it does not predict the {\\em scattered-light} component of the interstellar ultraviolet radiation field.  It is that defect that I repair in the present paper.  There is another reason for wanting to know the brightness and spectrum of the interstellar radiation field in the ultraviolet, quite apart from its influence on interstellar chemistry.  Henry (1991, 1999) has claimed that the spectrum of diffuse ultraviolet background radiation that is observed at high galactic latitudes has a sharp break near 1216 \\AA, with only an upper limit measurable below that wavelength, and a roughly constant spectrum at 300 to 400 \\phunit longward of that wavelength, to 2800 \\AA\\ (as appears in the spectrum of Anderson et al. 1979a,b), and continuing through the visible portion of the spectrum (Bernstein 1998; Bernstein, Freedman, and Madore 2002, ApJ, in press), and he has suggested that this radiation originates in recombination radiation from the baryonic component of the intergalactic medium.  A much simpler and more conservative interpretation is that the light in the ultraviolet at high galactic latitudes is simply galactic starlight that is reflected by high latitude interstellar dust:  some people have argued that there {\\em is} plenty of such dust (perhaps $A_V=0.1$, according to the assessment of Henry and Murthy (1994)) at high galactic latitudes (as revealed, for example, by IRAS) and if that is so, then the observed decline at the shortest wavelengths {\\em might} be due to nothing more than a decline at short wavelengths in the intensity of the general interstellar radiation field.  (If that view, that there is plenty of dust, is wrong, and there is in fact little dust at high galactic latitudes, then of course the alternative explanation collapses immediately.)  Henry's argument against this simpler picture has been his citation of the observation of the Coalsack nebula by Murthy, Henry, and Holberg (1994), who found that the bright ultraviolet spectrum of reflecting interstellar dust toward the Coalsack, which they observed with {\\em Voyager}, most definitely does {\\em not} decline below Lyman $\\alpha$.  The difficulty with this argument is that the Coalsack is being illuminated primarily by three stars ($\\alpha$~Cru, B0.5~IV, $V=1.58$; $\\beta$~Cru, B0.5~III, $V=1.25$; and $\\beta$~Cen, B1~III, $V=0.61$) that are among the brightest ultraviolet stars in the sky, and, specifically, that these particular stars are exceptionally hot and are also very close physically to the dust that they are illuminating, so that their light can be expected to dominate the scattered radiation. Thus, one could argue that, in  this case, the illumination is atypical, and that the general interstellar radiation field (which is what actually illuminates the high-galactic-latitude dust) has a cooler spectrum that does indeed decline sharply below Lyman $\\alpha$.   Only by determining the actual expected spectrum of the interstellar radiation field can one rule out this possibility, and a desire to assess this (yea or nay) formed a major motivation for the present work. ",
        "conclusions": "In contrast to my first exploration (Henry 1977) of the use of star catalog integrations as a means of predicting the interstellar radiation field, I have not presented (or obtained) maps of the predicted distribution of the radiation on the sky.  Any such predictions require knowledge of the shape of the scattering function of the interstellar grains, and are incomparably more uncertain than is the evaluation of the integrated spectrum of the radiation field as a whole, on which I have here exclusively focused.  I conclude that the albedo of the interstellar grains in the far ultraviolet is very low, perhaps $a=0.1$.  I emphasize that the present determination is much better founded than any previous determination by other methods.  The only significant uncertainty is the absolute calibration of the {\\em Apollo 17} ultraviolet spectrometer, but I have great confidence in the accuracy of that calibration.  In the present investigation I can come to no independent conclusion concerning the value in the far ultraviolet of the Henyey-Greenstein scattering parameter $g$, but I am happy to accept the conclusion of Witt et al. (1992) that the value of $g$ is large, corresponding to very strong forward scattering.  Burgh (2001) reaches the same conclusion.  Finally, I am able to conclude that Henry (1991, 1999) is correct in asserting that virtually all of the ultraviolet background radiation that is seen at high (and even at moderate) galactic latitudes is extragalactic in origin."
    },
    "0201/astro-ph0201205_arXiv.txt": {
        "abstract": "SDAMS is the ensemble of database + software packages aimed to the archiving, quick-look analysis, off-line analysis, network accessibility and plotting of the SPOrt produced data. Many of the aspects related to data archiving, analysis and distribution are common to almost all the astronomical experiments. SDAMS ambition is to face and solve problems like accessibility and portability of the data on any hardware/software platform in a way as simpler as possible, though effective. The system is conceived in a way to be used either by the scientific community interested in background radiation studies or by a wider public with low or null knowledge of the subject. The user authentication system allows us to apply different levels of access, analysis and data retrieving. SDAMS will be accessible through any Web browser though the most efficient way to use it is by writing simple programs. Graphics and images useful for outreach purposes will be produced and put on the Web on a regular basis. ",
        "introduction": "SPOrt (\\cite{carretti:2002}) on ISS is an experiment which will produce a limited amount of data with respect to other space experiments. The foreseen total (i.e. housekeepings and scientific data) bit rate is $\\sim 2.5$ kbit s$^{-1}$. This means that in 3 years of operation time it will collect $\\sim 30$ Gbytes of data. In addition SPOrt has not pointing capabilities and it is expected to work in a stable and semi-automatic way. As a consequence the output data stream will be simply identified by {\\em time} and {\\em orbital parameters} of the Space Station. The data management system (i.e. quick look, archive and scientific analysis) can be build to be both simple and highly automatic.  The SDAMS system has two fundamental parts: {\\em database server} and {\\em application server}. The former performs the data storing and allows a simple (low level) fast access to them. To this aim we chose MySQL$^{\\rm TM}$, a very easy to use data management system capable to deal with a relatively large amount of data. On the Web there are several examples of intensive usage of MySQL with application tools and there are extensions to use it within several programming languages (e.g. PHP). It has also the advantage to be an Open Source package, that means one can modify the source code to match any particular need and no license is required to use it. This is crucial to allow a standard (system) to be developed.  The SDAMS application server is a multithreaded server which uses a simple C-written communication protocol to interact with its clients. A task-oriented module library allows us to easily add further capabilities to those so far foreseen. The fundamental task of this component is to accept client connections (through socket) and invoke the appropriate modules  to perform the required tasks.  News about SDAMS will be posted on \\url{http://sport.tesre.bo.cnr.it}. ",
        "conclusions": ""
    },
    "0201/astro-ph0201496_arXiv.txt": {
        "abstract": "The 2.12\\micron\\ v=1--0 S(1) line of molecular hydrogen has been imaged in the Hourglass region of M8\\@. The line is emitted from a roughly bipolar region, centred around the O7 star Herschel 36.  The peak H$_2$ 1--0 S(1) line intensity is $\\rm 8.2 \\times 10^{-15} \\ erg \\ s^{-1} \\ cm^{-2} \\ arcsec^{-2}$.  The line centre emission velocity varies from $-25$\\,\\kms\\ in the SE lobe to $+45$\\kms\\ in the NW lobe. The distribution is similar to that of the CO J=3--2 line.  The H$_2$ line appears to be shock-excited when a bipolar outflow from Herschel 36 interacts with the ambient molecular cloud.  The total luminosity of all H$_2$ lines is estimated to be $\\rm \\sim 16 \\ L_{\\odot}$ and the mass of the hot molecular gas $\\rm \\sim 9 \\times 10^{-4} \\ M_{\\odot}$ (without any correction for extinction). ",
        "introduction": "\\label{sect:intro} The Lagoon Nebula, Messier 8, is an H{\\small II} region centred on the stellar cluster NGC 6523. It is embedded within a molecular cloud which extends to the young ($\\sim 2 \\times 10^6$ years old) star cluster NGC\\, 6530, $10'$ to its east (Lada et al.\\ 1976). Star formation is presumed to have proceeded from NGC\\,6530 and is now active in M8. Within M8's core lies the O7{\\small V} star Herschel 36 (Woolf 1961), which has created a blister-type H{\\small II} region, the Hourglass.  This visually distinctive nebula is extended $15''$ EW and $30''$ NS, and lies $15''$ E of Herschel 36.  The Hourglass is embedded within an extended H{\\small II} region, $\\sim 3'$ in extent, which is ionized by two O stars, HD\\,165052 and 9\\,Sgr (Woolf 1961, Lada et al.\\ 1976).  Nearby to Herschel 36 are a number of obscured sources, first observed in the near--IR by Allen (1986). Woodward et al.\\ (1990) have designated them KS1 to KS5\\@. Together, they may form a cluster of hot stars, analogous to the Trapezium in the Orion Nebula (M42), where $\\rm \\theta^1_c \\ Ori$ is the dominant member. In the mid--IR MSX\\footnote{see http://irsa.ipac.caltech.edu/ for details.} sky survey, this region appears as bright, extended source, with a 21\\micron\\ continuum flux of 960\\,Jy. This source is also prominent in sub--mm (Tothill 1999) and mm (Richter, Stecklum \\& Launhardt 1998) wavelength continuum. Nearby to Herschel 36, narrow band optical imaging with the HST reveals an object reminiscent of the proplyds in the Orion Nebula---a star within a bow-shock arc, whose apex is pointed towards Herschel 36 (Stecklum et al.\\ 1998). It is presumed to be an externally ionized circumstellar disk.  These authors also estimate the distance to M8 to be 1.8\\,kpc, based on the association with NGC\\,6530.  Of particular interest to this paper are the observations by White et al.\\ (1997) of intense CO line emission from M8.  They found the peak CO J=3--2 intensity to be over 100\\,K, making it the second brightest CO line source known. The CO J=3--2 and 4--3 lines were mapped and found to have a loose, bipolar structure, extending NW--SE from Herschel 36.  Taking the broad CO line profiles (extending over 20 \\kms), and the presence of a jet-like object extending $0.5''$ SE of Herschel 36 in HST images (Stecklum et al.\\ 1995), it seems likely that there is a molecular outflow in the core of M8\\@.  If so, it would be expected for there to be molecular hydrogen line emission as well, both shocked (from the deceleration of the outflow by the ambient cloud) and fluorescent (excited by far--UV photons from Herschel 36).  However, White et al.\\ (1997) also searched for the near--IR H$_2$ v=1--0 S(1) line in M8, but did not detect it.  In this paper we report more sensitive observations for H$_2$ emission from M8, and find that there is indeed excited H$_2$ line emission from this source. ",
        "conclusions": "Strong near--IR molecular hydrogen line emission is produced from around the Hourglass in the M8 star forming region.  The peak H$_2$ brightness is at the detection limit of a previous attempt to measure H$_2$ line emission from the source, thus explaining the reported non-detection.  The H$_2$ is emitted from an extended, roughly bipolar region, centred on the powering source for the Hourglass, Herschel 36. It is orientated along a NW--SE direction.  The morphology is similar to that of the CO J=3--2 distribution in M8.  Taken with the $\\sim 70$\\,\\kms\\ variation in the H$_2$ line centre across the region, and the 20\\,\\kms\\ width of the CO profiles, this suggests that the H$_2$ emission is shock-excited, when a bipolar outflow, originating from Herschel 36, impacts the surrounding molecular cloud.  The total H$_2$ line luminosity is $\\sim \\rm 16 \\ L_{\\odot}$ (not corrected for any extinction).  The H$_2$ lines provide a significant means of ridding the source of the mechanical energy in the outflow, though we cannot yet determine whether H$_2$ line emission is the dominant coolant in the shocked gas.  To examine the excitation mechanism more closely, in particular to determine whether there may also be fluorescent H$_2$ emission, it would be necessary to measure higher excitation lines of the molecule, from $\\rm v \\ge 2$.  This could be done either by long-slit spectroscopy, or through further Fabry-P\\'{e}rot imaging, with the waveplates tuned to appropriate lines.  To improve the determination of the flow energetics it will be necessary to use greater spectral resolution and resolve the line profiles, in order to measure the momentum and kinetic energy distribution of the hot molecular gas across the source."
    },
    "0201/astro-ph0201175_arXiv.txt": {
        "abstract": "We report here the results of a 90 ks BeppoSAX observation of the low mass X-ray binary and atoll source KS 1731--260 during a quiescent phase. From this observation we derive a source X-ray luminosity of $\\sim 10^{33}$ ergs/s (for a source distance of 7 kpc).  If the neutron star is spinning at a period of a few milliseconds, as inferred from the nearly-coherent oscillations observed during type-I X-ray bursts, the quiescent X-ray luminosity constrains the neutron star magnetic field strength.  We consider all the mechanisms that have been proposed to explain the quiescent X-ray emission of neutron star X-ray transients and compare the corresponding expectations with the measured upper limit on the X-ray luminosity. We find that, in any case, the neutron star magnetic field is most probably less than $\\sim 10^{9}$ Gauss. We have also observed KS 1731--260, still in its quiescence state, at 1.4 GHz with the Parkes radiotelescope to search for radio pulses. We found that no radio signals with millisecond periods are present with an upper limit on the flux of 0.60 mJy using a 4 min integration time (optimal for a close system with an orbital period smaller than a few hours) and of 0.21 mJy using a 35 min integration time (optimal for a wide orbit system). ",
        "introduction": "The Galactic X--ray source KS~1731--260 was discovered in August 1989 with the imaging spectrometer TTM on the Mir--Kvant observatory (Sunyaev 1989; Sunyaev et al. 1990).  During the $\\sim 15$~days long Mir--Kvant observation the source intensity varied from 50 to 100 mCrab in the 2--27~keV band. The presence of three type--I bursts indicated that the compact object is an accreting neutron star (hereafter NS) and that the source distance is about 7 kpc (Muno et al. 2000). The factor of 10 intensity variations displayed by the source suggested that  KS~1731--260 is a transient source (e.g.\\ Sunyaev et al. 1990). There were numerous detections of the source at a level of 50--100 mCrab since its discovery in 1989, and the monitoring by the Rossi X-ray Timing Explorer (RXTE) All Sky Monitor (ASM) showed that the source was continuously active until February 2001. Therefore the source appeared to belong to a group of NS soft X-ray transients (hereafter SXTs) that display active states extending for a number of years (as opposed to outbursts decaying on timescales of a few months; see Campana et al. 1998b for a review).  The source spectrum resembles that of other SXTs in outburst; the TTM spectrum taken on August 16--31, 1989 could be fit to a thermal bremsstrahlung with $kT \\sim 5.7$ keV while the spectrum obtained by SIGMA on  March 14, 1991 (when the source was somewhat fainter, $9 \\times 10^{36}$ erg/s as opposed to $1.5 \\times 10^{37}$ erg/s) was well modelled by a power law with a photon index of $\\sim 2.9$ extending up to 150 keV at least. The source detection during the ROSAT all-sky survey yielded a fairly accurate measurement of column density, $N_H = (1.00 \\pm 0.19) \\times 10^{22}$ cm$^{-2}$ (Barret, Motch, \\& Predehl 1998).  Based on an RXTE/PCA observation in July 1996, Smith, Morgan, \\& Bradt (1997) discovered a nearly coherent signal at $523.92 \\pm 0.05$ Hz (corresponding to a period of 1.91 ms) during a two-second interval close to the beginning of the decay of a type I X-ray burst. This signal, as well as similar signals observed during type I X-ray bursts from about ten low mass X-ray binaries, likely corresponds to the NS rotation frequency (or twice its value; for a review see Strohmayer, Swank, \\& Zhang 1998). An RXTE observation in August 1, 1996 (Wijnands \\& van der Klis 1997) revealed two simultaneous quasi periodic oscillations (QPO) in the persistent emission of KS 1731--260 at 898 Hz and 1159 Hz, the frequency separation of which ($260.3 \\pm 9.6$ Hz) is compatible with half the frequency of the nearly-coherent signal observed in the type--I bursts. In the beat frequency model for the kHz QPOs, this corresponds to a NS spinning at a period of $3.8$ ms, which might imply that only the first harmonic is detected during type--I bursts (although this case seems to be unlikely, see Muno et al. 2000).  RXTE/ASM data showed that the source entered a quiescent state in February 2001 (upper limit of 2 count/s corresponding to a luminosity of $< 3 \\times 10^{36}$ ergs/s). A few months after (on 27 March 2001) KS 1731--260 was observed by Chandra: it was still in quiescence with a 0.5--10 keV bolometric luminosity of $\\sim 10^{33}$ ergs/s (assuming a distance of 7 kpc) and an X-ray spectrum that is well described by a blackbody at a temperature of $\\sim 0.3$ keV (Wijnands et al. 2001a). The same data were also analyzed by Rutledge et al. (2001); they fitted the spectrum using a hydrogen atmosphere model, obtaining an effective temperature of 0.12 keV, an emission area radius of $\\sim 10$ km and a bolometric luminosity of $\\sim 2.7 \\times 10^{33}$ ergs/s (assuming a distance of 8 kpc). We present here the results of a 200 ks BeppoSAX observation of KS 1731--260 taken in March 2001, just after the beginning of this quiescent state, in which the source was detected at a quiescence luminosity level of $(7 \\pm 2) \\times 10^{32}$ ergs/s. ",
        "conclusions": "The 90 ks BeppoSAX observation of KS 1731--260 carried out on 2001 March 2--5 led to a detection of the source at a luminosity level of $\\sim 10^{33}$ ergs/s. This is similar to the flux determined in a subsequent observation of the source carried out with Chandra on 2001 March 27 (Wijnands et al. 2001a). These results testify that the quiescent X-ray luminosity and spectrum of KS 1731--260 are close to those determined for other NS SXTs. We therefore assume that the X-ray spectrum of KS 1731--260 in quiescence is similar to that of Aql~X-1 and Cen~X-4 (see e.g. Campana et al. 1998a), the best studied cases. The spectrum of these two sources could be well fit by a soft thermal component (blackbody temperature of $\\sim 0.2$~keV) plus a power-law component with a photon index $\\Gamma \\sim 1.5$. The blackbody component is usually interpreted as thermal emission from a pure hydrogen NS atmosphere (e.g.\\ Rutledge et al. 1999, 2000), while the power-law component is thought to be due to residual accretion or the interaction of a pulsar wind with matter released by the companion star (see e.g.\\ Campana \\& Stella 2000 and references therein).  In the following we discuss the properties of KS 1731--260 in relation to the different mechanisms that have been proposed to explain the quiescent X-ray emission of NS SXTs. In summary, there exist three sources of energy which are expected to produce X-ray luminosity in quiescence: \\begin{itemize} \\item[a)] residual accretion onto the NS surface at very low rates (e.g.\\ Stella et al. 1994); \\item[b)] rotational energy of the NS converted into radiation through the emission from a rotating magnetic dipole, a fraction of which can be emitted in X-rays (e.g.\\ Possenti et al. 2001; Campana et al. 1998b, and references therein); \\item[c)] thermal energy, stored into the NS during previous phases of accretion, released during quiescence (e.g.\\ Brown, Bildsten, \\& Rutledge 1998; Colpi et al. 2001; Rutledge et al. 2001). \\end{itemize} We will compare the corresponding expectations with the measured upper limit on the X-ray luminosity. Note that the last process has to be always taken into account, as the source was recently active for more than one year at a luminosity of $\\sim 10^{37}$ erg/s and frequently active, at comparable luminosities, since its discovery in 1989. Therefore we firstly derive some constraints resulting from the processes a) and b), and then we combine these with the constraint derived from process c).  Some constraints for the processes a) and b) can be derived if we assume that the coherent pulsations detected during the type I X-ray burst on 1996 July 14 (Smith et al. 1997) correspond to the NS spin.  Indeed the high degree of coherence ($\\nu / \\Delta \\nu \\ga 900$) is compatible with broadening induced by the finite duration of the oscillations ($\\sim 2$ s). This strongly supports the interpretation of the frequency of these oscillations ($\\nu = 523.92$ Hz) as the rotational frequency of the NS. Note that the frequency separation between the kHz QPOs in KS 1731--260 is $\\sim 260$ Hz (Wijnands \\& van der Klis 1997), which might be interpreted as the spin frequency of the NS (Strohmayer et al. 1996; Miller, Lamb, \\& Psaltis 1998), while the burst frequency would be its first overtone in this case (see, however, Muno et al. 2000). For concreteness in the following we will assume the burst frequency as the NS spin frequency and will discuss differences when needed.  It has been demonstrated that a rotating magnetic dipole in vacuum emits electromagnetic dipole radiation. Moreover a wind of relativistic particles associated with magnetospheric currents along the field lines is expected to arise in a rotating NS (e.g.\\ Goldreich \\& Julian 1969). Both these processes, powered by the rotational energy of the NS, depend on the angle between the NS magnetic moment and the spin axis and compensate in such a way that the total energy emitted is nearly independent of this angle (Bhattacharya \\& van den Heuvel 1991). Thus the bolometric luminosity of a rotating NS in vacuum can be calculated according to the Larmor's formula $L_{\\rm bol}= (2/3c^3)\\mu^2 \\omega^4$, where $c$ is the speed of light, $\\mu = B_{\\rm s} R_{\\rm NS}^3$ (where $B_{\\rm s}$ is the surface magnetic field at the magnetic equator and $R_{\\rm NS}$ is the NS radius), and $\\omega$ is the angular frequency of the NS. A small fraction of this luminosity is emitted in the radio band, probably as the results of electron accelerations in ultra-strong electric potentials (gaps). A still open question is the location of the accelerator (gap), in the outer magnetosphere, close to the light cylinder radius ($r_{\\rm LC} = c P / 2 \\pi$, i.e.\\ the radius at which an object corotating with the NS attains the speed of light), as proposed for the pulsar emission mechanism by Halpern \\& Ruderman (1993), or close to the magnetic cap, as originally proposed by Arons (1981) and Daugherty \\& Harding (1982). The observability radio emission can be strongly affected by the matter surrounding the NS. In particular free-free absorption can strongly reduce the radio flux especially at low radio frequencies (see e.g.\\ Burderi \\& King 1994, and, more recently, D'Amico et al. 2001). However, independent of the observability of the radio pulsar, this radiative regime certainly occurs when the space surrounding the NS is free of matter up to $r_{\\rm LC}$ and the pressure of this radiation can overcome the pressure of the accretion flow, thus preventing further accretion (see Ruderman, Shaham, \\& Tavani 1989; Illarionov \\& Sunyaev 1975). Therefore, in the following, we will adopt the hypothesis that once the magnetospheric radius is outside the light cylinder radius (and thus the space surrounding the NS up to the light cylinder is free of matter) the NS emits radiation according to the Larmor's formula. In this case, irrespective of the amount of radiation emitted in the radio band and of its observability, we will say, for shortness, that the radio pulsar is active.  As mentioned above, there are two possibilities: the first possibility, scenario a), is that the magnetospheric radius is inside the light cylinder radius and therefore the radio pulsar is off. In this case, for a non-zero NS magnetic field, we should have some matter flow toward the NS in order to keep the magnetospheric radius small enough. In this scenario we again have two possibilities: a1) the magnetospheric radius is inside the co-rotation radius (see below for a definition), so accretion onto the NS surface is possible; a2) the magnetospheric radius is outside the co-rotation radius (but still inside the light cylinder radius), so the accretion onto the NS is centrifugally inhibited, but an accretion disk can still be present and emit X-rays. The other possibility, scenario b), is that the magnetospheric radius is outside the light cylinder radius and therefore the radio pulsar is on. In this case X-ray emission can be produced by: b1) reprocessing of part of the bolometric luminosity of a rotating NS (Larmor's radiation formula) into X-rays in a shock front; b2) the intrinsic X-ray emission of the radio pulsar.  Let us consider scenario a), i.e.\\ a non-zero accretion rate during the quiescent phase, and in particular scenario a1). If the NS has an intrinsic magnetic field, we can estimate an upper limit on the NS magnetic moment. Note that a similar argument was applied to an ASCA observation of the Rapid Burster in quiescence by Asai et al. (1996). Also in that case it was concluded that a highly magnetized and rapidly rotating NS can be excluded. The matter accreting onto the NS forms an accretion disk whose inner radius is truncated at the magnetosphere by the interaction of the accretion flow with the magnetic field of the NS. In this case the magnetospheric radius $r_{\\rm m}$ is a fraction $\\phi \\la 1$ (an expression for $\\phi$ can be found in Burderi et al. 1998\\footnote{$\\phi = 0.21 \\alpha^{4/15} n_{0.615}^{8/27} m^{-142/945} [(L_{37}/\\epsilon)^{8/7} R_6^{8/7} \\mu_{26}^{5/7}]^{4/135}$, see the text for the definition of the symbols.}; for $L \\sim 10^{33}$ ergs/s we get $\\phi \\sim 0.2$) of the Alfv\\'en radius $R_{\\rm A}$ defined as the radius at which the energy density of the (assumed dipolar) NS magnetic field equals the kinetic energy density of the spherically accreting (free falling) matter: \\begin{equation} R_{\\rm A} = 2.23 \\times 10^6  R_6^{-2/7}  m^{1/7} \\mu_{26}^{4/7} \\epsilon^{2/7} L_{37}^{-2/7} \\;{\\rm cm} \\end{equation} (see, e.g., Hayakawa 1985), where $R_6$ is the NS radius, $R_{\\rm NS}$, in units of $10^6$ cm, $m$ is the NS mass in solar masses, $\\mu_{26}$ is the NS magnetic moment in units of $10^{26}$ G cm$^3$, $\\epsilon$ is the ratio between the specific luminosity and the specific binding energy ($L = \\epsilon \\times G M \\dot{M}/R_{\\rm NS}$, $G$ is the gravitational constant, $M$ is the NS mass, $\\dot{M}$ is the accretion rate), and $L_{37}$ is the accretion luminosity in units of $10^{37}$ erg/s, respectively. Actually, accretion onto a spinning, magnetized NS is centrifugally inhibited once the magnetospheric radius is outside the corotation radius, the radius at which the Keplerian frequency of the orbiting matter is equal to the NS spin frequency: \\begin{equation} r_{\\rm co} = 1.5 \\times 10^6 P_{-3}^{2/3} m^{1/3} \\;{\\rm cm} \\end{equation} where $P_{-3}$ is the spin period in ms. The condition to allow accretion then reads $r_{\\rm m}/r_{\\rm co} \\le 1$. This requires: \\begin{equation} \\mu_{26} \\le 0.5 \\phi^{-7/4} R_6^{1/2} \\epsilon^{-1/2} L_{37}^{1/2} m^{1/3} P_{-3}^{7/6}. \\end{equation} Adopting our upper limit, $L_{37} = 10^{-4}$, for the quiescent luminosity we obtain $\\mu_{26} \\le 0.28 m^{1/3}$ for $P_{-3} = 1.91$ ($\\mu_{26} \\le 0.63 m^{1/3}$ for $P_{-3} = 3.82$), where we have assumed $\\epsilon = R_6 = 1$ and $\\phi \\simeq 0.2$.  If the magnetospheric radius is outside the corotation radius and still inside the light cylinder radius (scenario a2) the accretion is centrifugally inhibited, but the disk can still emit X-rays. This will increase the upper limit on the magnetic field derived above because in the propeller regime the system will be under-luminous, by a factor $\\epsilon_{\\rm prop} = R_{\\rm NS}/(2 r_{\\rm disk})$, for a given accretion rate. In this case, the maximum accretion rate for a given luminosity occurs when the efficiency factor $\\epsilon$ is minimum.  This occurs when the inner disk radius is the furthest from the NS surface without allowing the switch-on of the radio pulsar, i.e. for $r_{\\rm disk} \\la r_{\\rm LC}$. Assuming $r_{\\rm m} = r_{\\rm LC}$ and taking as upper limit on $\\dot M$ the condition $L_q = L_{\\rm disk} (r_{\\rm LC}) = \\epsilon_{\\rm prop} G M \\dot{M}/R_{\\rm NS}$, with $\\epsilon_{\\rm prop} = R_{\\rm NS}/(2 r_{\\rm LC})$, we can calculate the upper limit on the magnetic moment in this case: \\begin{equation} \\mu_{26} \\le 11.7 \\phi^{-7/4}  L_{37}^{1/2}  P_{-3}^{9/4} m^{-1/4}, \\end{equation} where we have assumed $R_6 = 1$. This gives $\\mu_{26} \\le 8.4 m^{-1/4}$ for $P_{-3} = 1.91$ (and $\\mu_{26} \\le 39.8 m^{-1/4}$ for $P_{-3} = 3.82$), where we have assumed $L_{37} = 10^{-4}$ for the quiescent luminosity, and $\\phi \\simeq 0.2$ (in fact, as $\\phi \\propto \\epsilon_{\\rm prop}^{-32/945}$, $\\phi = 0.18 \\simeq 0.2$ is also appropriate in the propeller case).  Let us consider now the scenario b). If there is no accretion in quiescence or the magnetospheric radius falls outside the light cylinder radius, we expect that the radio pulsar switches on. The pulsar radiation pressure may overcome the pressure of the accretion disk, thus determining the destruction of the disk and the ejection of matter from the system (see Burderi et al. 2001). However there is the possibility that a fraction $\\eta$ of this power can be converted into X-rays in a shock front between the emerging radiation and the circumstellar matter (scenario b1). Typical values for $\\eta$ are in the range $\\sim 0.01-0.1$ (e.g.\\ Campana et al. 1998b). In this case $\\eta L_{\\rm rad} \\la L_q$ or \\begin{equation} \\mu_{26} \\le 5.1 L_{37}^{1/2} P_{-3}^2 \\eta^{-1/2}. \\end{equation} In our case, assuming the minimum value for the efficiency in the conversion of the rotational energy into X-rays, $\\eta = 0.01$, which implies the highest value for the magnetic moment, we get $\\mu_{26} \\le 1.9$ for $P_{-3} = 1.91$ (and $\\mu_{26} \\le 7.5$ for $P_{-3} = 3.82$). Let us consider the possibility that part of the spin-down energy loss is directly emitted in X-rays (scenario b2). Indeed about 40 out of $\\sim 1400$ known radio pulsars have been detected in the X-ray range so far. In this case a correlation has been observed between the X-ray and spin-down luminosities (e.g.\\ Becker \\& Trumper 1997; Possenti et al. 2001, and references therein).  Considering the empirical relation, derived by Possenti et al. (2001) analyzing a sample of 37 pulsars, between the 2--10 keV luminosity, $L_{37}$, and the rate of spin-down energy loss, $L_{\\rm rad}$, \\begin{equation} L_{37} = 2.51 \\times 10^{-52} L_{\\rm rad}^{1.31}, \\end{equation} we can again calculate an upper limit on the magnetic moment in the case of KS 1731--260, which gives $\\mu_{26} \\le 8.7$ for $P_{-3} = 1.91$ (and $\\mu_{26} \\le 34.6$ for $P_{-3} = 3.82$).  In the last two cases, the X-ray emission is expected to be non-thermal (a power-law spectrum could be appropriate). Interestingly Wijnands et al. (2001) found that the Chandra spectra of KS 1731--260 in quiescence could be equally well fitted by a blackbody or a power law plus a blackbody (although the power law is not statistically required). In this second case one might think that the thermal component arises from the NS surface and the non-thermal power law is associated with the shock emission discussed above. Indeed, according to the results of the spectral deconvolution of Wijnands et al. (2001), the fraction of the total luminosity emitted in the power law is $\\sim 15\\%$, small as compared to the total luminosity. However, because of the low statistics, the spectral deconvolution is not secure and, in the case of our BeppoSAX observation, no statistically significative spectral analysis is possible. For this reason, in deriving our upper limits on the magnetic moment strength in these scenarios, we have adopted the very conservative assumption that the whole flux detected by BeppoSAX can be ascribed to the non-thermal emission.  An upper limit on the NS magnetic field can also be derived from the upper limit on the radio pulsed emission from KS 1731--260. Radio observations of KS 1731--260, taken at the Parkes radiotelescope while the system was in quiescence, showed that the upper limit on the flux of pulsed radio emission is 0.21 mJy using 35 min integration time and 0.60 mJy using 4 min integration time.  Equating the rotational energy lost by a magnetized NS  with its magnetic dipole radiation, we obtain: \\begin{equation} P \\dot{P} = 9.75 \\times 10^{-24} \\mu_{26}^2 I_{45}^{-1} \\label{eq:ppdot} \\end{equation} where $P$ and $\\dot{P}$ are respectively the period of the radio pulsar and its derivate and $I_{45}$ is the NS momentum of inertia in units of $10^{45}$ g cm$^2$. Proszynski \\& Przyibician (1985), using a sample of 275 pulsars, found an empirical relation between the observed 400 MHz luminosity, $L_{400}$, and the pulsar parameters: \\begin{equation} \\log(L_{400})=\\frac{1}{3} \\log\\left(\\frac{\\dot{P}_{-15}}{P^3}\\right) + A \\label{eq:L400} \\end{equation} where $L_{400}$ is expressed in  mJy kpc$^2$, $\\dot{P}_{-15}$ is in units of $10^{-15}$ s s$^{-1}$ and $A  = 1.1$. Kulkarni,  Narayan  \\&  Romani (1990) found that the  same relation could easily fit a sample of 11 recycled binary pulsars. For KS 1731--260, assuming a distance of 7.0 kpc and a spectral index $\\alpha = 1.7$, typical of millisecond pulsars, we obtain for the 400 MHz luminosity an upper limit of 88.47 mJy kpc$^2$ in the case of 35 min integration, and of 250.25  mJy kpc$^2$ in the case of 4 min integration. Using  these values    and  combining  equations   \\ref{eq:ppdot}  and \\ref{eq:L400}, we can again calculate  an upper limit  to  the magnetic field of KS 1731--260. If the NS  is in a  wide orbit, i.e. if the phase shifts introduced  by the doppler effect are negligible over  an integration time  of 35 min,  we obtain  an  upper limit of $\\mu_{26} \\le 0.68$ for $P_{-3} = 1.91$ (and $\\mu_{26} \\le 2.7$ for $P_{-3} = 3.82$), assuming $I_{45}=R_6=1$. If, instead, the NS in KS 1731--260 belongs to a narrow binary  system and then the 4 min observations are  more appropriate, we obtain $\\mu_{26} \\le 3.1$ for $P_{-3} = 1.91$ (and $\\mu_{26} \\le 12.6$ for $P_{-3} = 3.82$). Note however that free-free absorption from matter surrounding the system (expected if the mass loss from the companion proceeds at some level during the quiescent phase, see Burderi et al. 2001 for a detailed discussion of this possibility in transient systems) can efficiently hamper the detection of the pulsed radio emission from an active radio pulsar.  We will consider now the process c), i.e.\\ the thermal emission from the NS. Deep crustal heating was first developed by Haensel \\& Zdunick (1990).  It was pointed out by Brown et al. (1998) that this scenario was relevant to accreting NSs in transient systems, and may contribute to the emission at low luminosities.  This scenario applies if the system has time to reach a steady state, in which the heat deposited during short and frequent outbursts in the NS is equal to the luminosity radiated in quiescence. Colpi et al. (2001) have explored the thermal evolution of a NS undergoing episodes of accretion, lasting for a time $t_{\\rm out}$, separated by periods of quiescence, lasting $t_{\\rm rec}$ in the hypothesis that the thermal luminosity in quiescence is dominated by emission from the hot NS core, with which the crust and atmosphere of the NS are in thermal equilibrium. Adopting an outburst luminosity of $10^{37}$ ergs/s and a quiescent luminosity of $10^{33}$ ergs/s, this model predicts a ratio of the recurrence time $t_{\\rm rec}$ to the outburst time $t_{\\rm out}$ of $\\sim 70$ for a 1.4 $M_\\odot$ NS and less than 4 for a NS more massive than 1.7 $M_\\odot$. A $t_{\\rm rec}/t_{\\rm out} \\sim 70$ would imply that, on average, the total time the source should spend in outburst in 13 yrs (i.e.\\ since its discovery) should be about 2 months. This is in conflict with time history of this source and the fact that the last outburst episode observed by the ASM lasted much more than one year. As already pointed out by Wijnands et al. (2001a), we are therefore left with two possibilities: i) we are witnessing a peculiar period of hyperactivity of the source, which will be followed by a much longer (up to a thousand years); ii) the NS in this system is quite massive, in agreement with the expectation of the standard conservative recycling scenario (e.g.\\ Bhattacharya \\& van den Heuvel 1991) that the NS has accreted a significant amount of mass and has been spun-up to a millisecond spin period.  However, Rutledge et al. (2001) found that, since KS 1731-260 has recently experienced an extremely long outburst, the thermal luminosity in quiescence could be dominated by emission from the heated NS crust out of thermal equilibrium with the core.  Indeed the prediction of the quiescent luminosity from the cooling crust model agrees (within a factor of a few) with the observed bolometric luminosity in quiescence. In this case the model above cannot be applied to the system, given that the core might be much cooler than indicated by the thermal luminosity in quiescence.  This implies that the presence of a massive NS is not required to explain the behavior of this system, and that the estimated long outburst recurrence timescale mentioned above should be considered as a lower limit. Rutledge et al. (2001) also calculated the expected thermal evolution of the crust, which should cool in 1--30 years, and therefore the time evolution of the quiescent luminosity for different crustal conductivities and cooling processes, showing that a monitoring of this source can provide valuable information on the crust and core microphysics.  A further constraint on the magnetic moment can be derived considering that the NS spin period must be longer than the equilibrium spin period, namely the Keplerian period at the magnetospheric radius: \\begin{equation} P_{\\rm eq} = 1.81 \\times 10^{-3} \\phi^{3/2} R_6^{-3/7} \\mu_{26}^{6/7} \\epsilon^{3/7} L_{37}^{-3/7} m^{-2/7} \\; {\\rm s.} \\end{equation} Adopting $\\phi = 0.2$, and an outburst luminosity of $L_{37} = 1$ (i.e.\\ the averaged luminosity observed with ASM and previous missions when the source was detected) we obtain $\\mu_{26} \\le 16.4 m^{1/3}$ for $P_{-3} = 1.91$ (and $\\mu_{26} \\le 36.8 m^{1/3}$ for $P_{-3} = 3.82$). As KS~1731--260 is a transient source it is possible that, during quiescence, the NS enters a propeller phase during which an efficient torque could spin down the NS far from the equilibrium period defined above. In this case the magnetic moment is significantly {\\it less} than the value derived in this paragraph. On the other hand, if the propeller phases are infrequent and/or the spin down torque is unefficient, the NS is expected to be in spin equilibrium and $\\mu_{26} = 16.4 m^{1/3}$ ($\\mu_{26} = 36.8 m^{1/3}$).  From the discussion above we conclude that, in any case, the NS magnetic field is most probably less than $\\sim 10^{9}$ Gauss, and less than $\\sim 4 \\times 10^{9}$ Gauss in the worst case.  Let us finally discuss the consequences of the spin equilibrium constraint when applied to the proposed scenarios in quiescence. In fact, if we consider the case that the NS in KS~1731--260 is quite massive, the amount of mass accreted, $\\ga 0.3$ $M_\\odot$, is sufficient, in principle, to spin up the NS below 1 ms for soft or moderately stiff equations of state for the NS matter and below 1.5 ms even for the stiffest equations of state (see Burderi et al. 1999). Therefore: \\begin{itemize} \\item[a1)] The NS in quiescence accretes matter onto its surface at very low rates. In this case $\\mu_{26} \\le 0.28 m^{1/3}$ ($\\mu_{26} \\le 0.63 m^{1/3}$). This, compared with the spin equilibrium constraint, i.e.\\ $\\mu_{26} \\le 16.4 m^{1/3}$ ($\\mu_{26} \\le 36.8 m^{1/3}$), implies that the NS is far from spinning at equilibrium, {\\it i.e.} propeller phases must be frequent and the spin down torque very effective, although this quiescent phase does not correspond to a propeller. \\item[a2)] The NS in quiescence is in a propeller phase. In this case $\\mu_{26} \\le 8.4 m^{-1/4}$ ($\\mu_{26} \\le 39.8 m^{-1/4}$). This, compared with the spin equilibrium constraint, $\\mu_{26} \\le 16.4 m^{1/3}$ ($\\mu_{26} \\le 36.8 m^{1/3}$), indicates that the NS is compatible with spinning close to the equilibrium. In this case propeller phases must be rare (although the present quiescent phase does correspond to a propeller) and/or the spin down torque not very effective. \\item[b1)] The NS in quiescence is not accreting matter and the radio pulsar is active; a fraction of the power emitted by the rotating NS is converted into X-rays in a shock front between the emerging radiation and the circumstellar matter. In this case $\\mu_{26} \\le 1.9$ ($\\mu_{26} \\le 7.5$). This, compared with the spin equilibrium constraint, implies that the NS is far from spinning at equilibrium, {\\it i.e.} again propeller phases must be frequent and the spin down torque very effective, although this quiescent phase does not correspond to a propeller. \\item[b2)] The NS in quiescence is not accreting matter and the radio pulsar is active; in this case the fraction of power emitted by the rotating NS that is converted into X-rays in a shock front between the emerging radiation and the circumstellar matter is negligible, and the X-ray emission is the intrinsic emission from the rotating NS. In this case $\\mu_{26} \\le 8.7$ ($\\mu_{26} \\le 34.6$). This implies that the NS is compatible with spinning close to the equilibrium, {\\it i.e.} propeller phases must be rare (indeed the present quiescent phase does not correspond to a propeller) and/or the spin down torque not very effective. \\end{itemize}  In conclusion, scenarios a2) and b2) seems to be the most reasonable, and both indicate that the NS is compatible with spinning close to the equilibrium period."
    },
    "0201/astro-ph0201343_arXiv.txt": {
        "abstract": "We have analysed arrival times of extensive air showers (EAS) registered with the \\EAS\\ prototype array during the period from August 1997 till February 1999. Our analysis has revealed that though the vast majority of samples of consecutive time intervals between EAS arrival times obey the exponential distribution, there are sequences of showers that have another distribution and thus violate the homogeneity hypothesis. The search for correlation between such events and clusters of showers and events with big delays between arrival times was also carried out. ",
        "introduction": "Distribution of arrival times of extensive air showers (EAS) has been studied by methods of classical statistics by different research groups (see, e.g., Chikawa et~al.\\ (1991) and Tsuji et~al.\\ (1993)).  In this paper, we present some of the results of a similar investigation performed with the experimental data obtained with the \\EAS\\ prototype array. This installation operates at the Institute of Nuclear Physics of Moscow State University (Fomin et~al., 1999).  The \\EAS\\ prototype array consists of eight detector units (DU) placed in the central part of the EAS MSU array and covers the area of $67\\times22$~m.  The array is designed for registration of extensive air showers produced by cosmic rays with energy more than $10^{14}$~eV.  When an EAS arrives, the corresponding signals from DU are transmitted to a PC for the preliminary program analysis of information quality and for EAS event verification.  The shower selection criterion is the triggering of any four adjacent detectors in the course of time gate less than 3.2~ns. ",
        "conclusions": "It should be noted that our results do not contradict the conclusions of other research groups (Chikawa et~al., 1991; Tsuji et~al., 1993).  However, it is necessary to increase shower statistics and to use other non-parametric methods of hypothesis verification since different criteria may give different results for the homogeneity assumption test (Bendat and Piersol, 1986)."
    },
    "0201/astro-ph0201425_arXiv.txt": {
        "abstract": "We have developed a physically self-consistent model of the disk around the nearby 10 Myr old star TW Hya which matches the observed spectral energy distribution and 7mm images of the disk.  The model requires both significant dust size evolution and a partially-evacuated inner disk region, as predicted by theories of planet formation. The outer disk, which extends to at least 140 AU in radius, is very optically thick at infrared wavelengths and quite massive ($\\sim 0.06 \\msun$) for the relatively advanced age of this T Tauri star.  This implies long viscous and dust evolution timescales, although dust must have grown to sizes of order $\\sim 1$~cm to explain the sub-mm and mm spectral slopes. In contrast, the negligible near-infrared excess emission of this system requires that the disk be optically thin inside $\\lesssim$ 4 AU.  This inner region cannot be completely evacuated; we need $\\sim $ 0.5 lunar mass of $\\sim$ 1 $\\mu$m particles remaining to produce the observed $10 \\mu$m silicate emission. Our model requires a distinct transition in disk properties at $\\sim 4$ AU, separating the inner and outer disk. The inner edge of the optically-thick outer disk must be heated almost frontally by the star to account for the excess flux at mid-infrared wavelengths. We speculate that this truncation of the outer disk may be the signpost of a developing gap due to the effects of a growing protoplanet; the gap is still presumably evolving because material still resides in it, as indicated by the silicate emission, the molecular hydrogen emission, and by the continued accretion onto the central star (albeit at a much lower rate than typical of younger T Tauri stars). TW Hya thus may become the Rosetta stone for our understanding of the evolution and dissipation of protoplanetary disks. ",
        "introduction": "The discovery of extrasolar planets (Marcy \\& Butler 1998 and references therein) has opened up a new era in the study of planetary systems.  While many important clues to the processes of planet formation can be obtained from studies of older systems, the best tests of formation scenarios will require the direct detection of actively planet-forming systems.  It is thought that the formation of giant planets involves the sweeping up of material in a wide annulus in the circumstellar disk, resulting in the development of a gap (Lin \\& Papaloizou 1986, 1993; Bryden \\etal 1999). Material inside the planet-driven gap can continue to accrete onto the central star; if the planet can prevent material from accreting across the gap into the inner disk, the eventual result would be the evacuation of the region interior to the planet.  In the case of the Solar System, the formation of Jupiter might have prevented outer disk gas from reaching the inner solar system; the inner gas disk accreted into the Sun, while solid planetesimals remaining behind eventually formed the terrestrial planets.  The above scenario suggests that the signature of a forming giant planet would be the presence of a gap that is not entirely evacuated. In this case dusty emission from the inner disk might still be observable. In addition, giant planet formation requires the prior consolidation of large solid  bodies to serve as cores for subsequent gas accretion; while such bodies would be invisible with current techniques, one would expect to see evidence for substantial growth in dust particles.  Finally, if the inner disk has not been completely evacuated by the forming planet, one might expect to observe continued accretion onto the central star, as all T Tauri systems with inner disks as detected from near-infrared disk emission are also accreting (Hartigan \\etal 1990).  In this article we propose that the relatively young low-mass star TW Hya has a developing gap in its inner disk qualitatively similar to that expected from planet formation. The evidence supporting this this proposal is: (1) reduced (optically-thin) emission from the inner disk; (2) mm-wave spectra which seem to require grain growth; (3) extra emission from the edge of the outer disk; and (4) continued accretion onto the central star, albeit at a rate substantially lower than that observed from most T Tauri stars.  TW Hya has an age of 10 Myr and so according to current theories it is quite likely to be close to the epoch of planet formation.  It is also part of an association of young stars of similar age, but it stands out in that it is the only system still accreting at a substantial rate (Muzerolle \\etal 2000).  The (outer) disk of TW Hya is quite massive, so that there is likely to be more than enough material available to form giant planet(s).  TW Hya is also the closest known such system, and so it will be a prime target for following studies to confirm our model. ",
        "conclusions": "The disk of TW Hya seems to be in advanced state of dust evolution.  A planet or other large perturbing body may have already formed, opening a gap with outer edge around $\\sim$ 4 AU, and although there is still material in the inner disk, it is rapidly dissipating. In the outer disk, grains may be reaching the size necessary to start settling towards the midplane. This interpretation can consistently explain the large degree of activity of TW Hya, the lack of disk emission at near-IR wavelengths and the large fluxes beyond $\\sim$ 10 $\\mu$m, as well as the 7 mm images.  Evacuated regions in the disks are also known to result from the effects of companion stars (GG Tau), and in some cases disk accretion can still occur, especially if the binary has an eccentric orbit (DQ Tau).  Indeed, HD98800, another member of the TW Hya association, shows evidence for an inner disk hole; because this system is quadruple, it is quite likely that the inner disk regions have been evacuated by the companion stars. Although we cannot rule out such a possibility in TW Hya, the fact that accretion still occurs onto the central star, which is not the case in HD 98800, suggests that the companion body responsible for the gap is much less massive than a typical companion star, and thus much less effective in clearing out the disk.  Interferometric imaging should be attempted to constrain the mass of any companion object, expected to have a separation of order 2-3 AU (0.04 - 0.06 arcsec at 55 pc).   Acknowledgments. We gratefully acknowledge discussions with Edwin Bergin and Stephen Lepp. This work has been supported by NASA Origins of Solar Systems grants NAG5-9670 and  NAG5-9475. P.D. has been supported by Conacyt grant J27748E and a fellowship from DGAPA-UNAM, M\\'exico.     \\newpage"
    },
    "0201/astro-ph0201339_arXiv.txt": {
        "abstract": "We analysed {\\it Rossi X-ray Timing Explorer} (RXTE) data of the low-mass X-ray binary and atoll source 4U 1608-52 obtained on March 3, 1996 in which the source simultaneously showed a strong single kilohertz quasi-periodic oscillation (QPO) around 840 Hz, and a 7.5 mHz QPO detected at energies below 5 keV.  We find that the frequency of the kHz QPO is approximately anti-correlated with the 2--5 keV X-ray count rate associated with the mHz QPO. The average kHz QPO frequency varies by about 0.6 Hz (0.07\\%) during a mHz QPO cycle over which the average 2--5 keV count rate varies by about 60 c/s (4\\%). This is opposite to the frequency-count rate correlation observed in the same data on longer time scales and hence constitutes the first example of a sign reversal in the frequency-flux correlation related to the origin of the flux. Such a sign reversal is predicted by the radiative disk truncation model for the case where the flux variations originate on the neutron star but are not due to disk accretion rate fluctuations. The results support the nuclear burning interpretation of the mHz QPO, and the interpretation of the kHz QPO frequency as an indicator of the orbital frequency at the inner edge of the accretion disk. The varying radiative stresses on the inner disk exerted by the flux due to the quasi-periodic nuclear burning lead to changes in the inner disk radius and hence to the observed anti-correlation between kHz QPO frequency and X-ray count rate. There is a time lag of about ten seconds of the X-ray count rate relative to the kHz QPO frequency in terms of the anti-correlation we found between these two quantities, which could be caused by the propagation of the nuclear burning front on the neutron star away from the equatorial region. ",
        "introduction": "Kilohertz quasi-periodic oscillations (kHz QPOs) in X-ray flux have been observed from about 20 neutron star low mass X-ray binaries (LMXBs) (see van der Klis 2000 for a recent review). Their frequencies correspond to the dynamical time scale at radii of less than a few tens of kilometers, and are usually regarded as associated with the Keplerian orbital frequency of material at the inner disk edge. Observations indicate that when on a time scale of hours to days the accretion rate through the disk increases, the QPO frequency increases as well, indicating that the inner disk edge moves in. Evidence for this has been observed in the form of a positive correlation on these time scales of kHz QPO frequency with X-ray count rate for low luminosity LMXBs, i.e. 'atoll' sources (see Mendez et al. 1999), and in the form of a similar correlation with curve length S$_z$ along the track traced out in an X-ray color-color diagram for high luminosity LMXBs, i.e. 'Z' sources, which in some cases corresponds to an anti-correlation between kHz QPO frequency and X-ray count rate (Wijnands et al. 1997; Homan et al. 2001; Yu, van der Klis \\& Jonker 2001).  The inner edge of the disk could be set by magnetosphere-disk interaction (Strohmayer et al. 1996, Zhang, Yu \\& Zhang 1998, Cui 2000, Campana 2000), and for sufficiently compact neutron stars with sufficiently weak magnetic field is certainly expected to be limited by the general relativistic marginally stable orbital radius. However, in the sonic point model for kHz QPOs (Miller, Lamb and Psaltis 1998, also Miller and Lamb 1995) radiative stresses provide the dominant disk truncation mechanism. In this description, removal of angular momentum from the flow by radiation drag, and the radial radiation force combine to set the inner disk radius and thereby the QPO frequency. For constant radiation, the disk moves {\\it in} when the disk mass flow increases; for constant mass flow, the disk moves {\\it out} when the radiation impinging upon the inner disk edge becomes more intense. Usually, disk mass flow and radiation change in concert; calculations show that under certain assumptions the inner disk radius will then move {\\it in} and hence QPO frequency will go {\\it up} when disk accretion and the resulting X-ray flux both increase, and {\\it vice versa} (Miller et al. 1998). If disk accretion is not the only mechanism producing radiation more complicated relations can occur (van der Klis 2001).  A test of the role of radiative stress for disk truncation can be obtained by studying the changes in kHz QPO frequency resulting from changes in X-ray flux that are {\\it not} due to disk mass flow changes. Two previous attempts at performing this test led to results that were somewhat ambiguous. (i) An anti-correlation was recently found between the kHz QPO frequency and the NBO (``normal branch oscillation'') count rate in Sco X-1 ($\\sim$22 Hz shift within the 1/6-s NBO cycle, Yu, van der Klis \\& Jonker 2000). However, Sco X-1 was in a state where such an anti-correlation is also seen for long-term flux changes that are presumably associated with disk flow changes. (ii) A 37-Hz decrease of the kHz QPO frequency was observed after an X-ray burst in Aquila X-1 (Zhang, Jahoda \\& Kelley et al. 1998; Yu, Li \\& Zhang et al. 1999). However, other similar bursts did not show this effect.  In this Letter, we show that in the atoll source 4U 1608--52 the kHz QPO frequency is clearly anti-correlated with the X-ray count rate variations associated with the 7.5 mHz QPO found by Revnivtsev, Churazov \\& Gilfanov et al. (2001). This QPO occurs in a narrow X-ray luminosity range and they suggested, among other possibilities, that these were due to a special mode of nuclear burning on the neutron star surface. This is the first clear-cut case of a sign reversal in the kHz QPO frequency vs. count rate correlation related to the basic process producing the radiation (disk accretion vs. nuclear burning). It provides strong evidence in favor of a role for radiative stresses in the disk truncation mechanism, as well as for the nuclear burning scenario for the mHz QPO. ",
        "conclusions": "We have studied the evolution of the lower kHz QPO as a function of the phase of the 7.5 mHz QPO displayed by 4U 1608--52 when its luminosity is about 0.1 $L_{Edd}$. The average kHz QPO frequency depends on the mHz QPO phase. It varies by about 0.6 Hz, or 0.07\\% of the average frequency, around the mHz QPO cycle, in rough anti-correlation to the average 2--5 keV count rate, which varies by about 60 c/s, or 4\\% of the average 2--5 keV count rate (or approximately 2\\% of the entire PCA 2--60 keV count rate). We have also found that in this anti-correlated variation the kHz QPO frequency leads the count rate by about ten seconds. On longer time scales (hours), the kHz QPO frequency in these same observations was already known to be positively correlated with the X-ray count rate and X-ray luminosity (Mendez et al. 1999), and our analysis confirms this.  The fact that we find an {\\it anti-correlation} on the 7.5 mHz QPO time scale, while simultaneous count rate changes not related to the mHz QPO produce the usual {\\it correlated} variations in kHz QPO frequency demonstrates conclusively that another mechanism than that responsible for the usual count rate variations underlies the mHz QPO.  The mHz QPO (Revnivtsev et al. 2001) occurs only within a narrow range of source luminosity and disappears after X-ray bursts. Its spectrum is rather soft and consistent with a blackbody. For these reasons, Revnivtsev et al. (2001) suggested that the mHz QPO could be due to a special mode of nuclear burning on the neutron star surface, although an origin due to processes in the accretion disk could not be ruled out. We suggest that the simplest explanation for the anti-correlation between kHz QPO frequency and mHz QPO count rate we report here is, that the mHz QPO is generated inside the inner disk edge by the mechanism proposed by Revnivtsev et al. (2001), and that the kHz QPO frequency is related to the Keplerian frequency at the inner edge of the disk, whose radius in turn is set by the radiative disk truncation mechanism proposed by Miller et al. (1998; see \\S1).  If instead the mHz QPO were produced by a modulation of the disk accretion rate, a positive correlation between kHz QPO frequency and mHz QPO X-ray flux would be predicted, contrary to what we observe. Indeed, such a positive correlation is seen for the longer-timescale flux variations.  If the mHz QPO indeed originates near the neutron star surface, i.e. inside the inner disk radius, an anti-correlation between the orbital frequency of disk material and the mHz QPO flux is a prediction of the radiative disk truncation mechanism, as the inner disk edge will move out when the X-ray flux impinging upon it increases, due to both an increased radiation drag and an increased radial radiation force (Miller et al. 1998).  Following their equation (39) for the effect of the radial radiation force, taking the luminosity of the neutron star as $\\sim$0.1 $L_{E}$, a 2\\% change in the neutron star luminosity will introduce a change in orbital frequency of about 0.1\\%, i.e. 0.8 Hz for an orbital frequency of 830 Hz, or 1.1 Hz for an orbital frequency of 1060 Hz for the case when the twin peak frequency separation is taken at a constant 230 Hz (see M\\'endez et al. 1998).  Radiation-drag effects will enhance the change in the kHz QPO frequency; also the change in neutron star luminosity will be larger than the 2\\% we adopted if not all the source flux originates from the neutron star surface (in most models, the hard flux originates elsewhere).  So, quantitatively the theoretical expectation is somewhat larger than the 0.6 Hz we observe. As suggested below, a more detailed analysis of the emission pattern caused by the nuclear burning may explain this.  The estimate of the frequency amplitude we made above was based on the assumption that the relative change in mHz QPO flux we observe on Earth is the same as that in the flux towards the inner disk edge. This can not be true if the mHz QPO is caused by local thermonuclear fires propagating on the neutron star. If the fires originate mostly in the equatorial region, where most of the accreted matter piles up, and then expand and propagate away, then the flux amplitude experienced by the disk would be less than that observed, and the radiative stresses exerted on the inner disk edge would reach their maximum before the observed X-ray flux does. This could explain both the lower than expected frequency amplitude and the time lag between the X-ray count rate and the kHz QPO frequency which we observed.  In this paper we have demonstrated that study of the correlation between low-frequency variability and high frequency QPOs can provide strong constraints on the origin of both variability modes. Further similar studies of the relations between low-frequency variability and, e.g., kHz QPOs in neutron star X-ray binaries and high frequency QPOs in black hole X-ray binaries, could potentially provide considerable additional insight into the nature of both QPOs and broad-band variability in these objects."
    },
    "0201/astro-ph0201049_arXiv.txt": {
        "abstract": "{ We present the first speckle interferometric observations of \\object{R CrB}, the prototype of a class of peculiar stars which undergo irregular declines in their visible light curves. The observations were carried out with the 6~m telescope at the Special Astrophysical Observatory near maximum light ($V=7$, 1996 Oct. 1) and at minimum light ($V=10.61$, 1999 Sep. 28). A spatial resolution of 75~mas was achieved in the $K$-band. The dust shell around R~CrB is partially resolved, and the visibility is approximately 0.8 at a spatial frequency of 10 cycles/arcsec. The two-dimensional power spectra obtained at both epochs do not show any significant deviation from circular symmetry. The visibility function and spectral energy distribution obtained near maximum light can be simultaneously fitted with a model consisting of the central star and an optically thin dust shell with density proportional to $r^{-2}$.  The inner boundary of the shell is found to be $82$~\\mbox{$R_{\\star}$}\\ (19~mas) with a temperature of $920$~K. However, this simple model fails to simultaneously reproduce the visibility and spectral energy distribution obtained at minimum light. We show that this discrepancy can be attributed to thermal emission from a newly formed dust cloud. ",
        "introduction": "The R Coronae Borealis (RCB) stars are a class of unusual objects characterized by sudden declines in their visible light curves as deep as $\\Delta V \\sim 8$. They are extremely hydrogen-deficient and also carbon-rich (e.g. Asplund et al. \\cite{asplund00} and references therein). The RCB stars are thought to undergo the formation of dust clouds in random directions, and it is believed that a sudden decline takes place, only when a dust cloud forms in the line of sight (Loreta \\cite{loreta34}, O'Keefe \\cite{okeefe39}). A newly formed dust cloud is expected to be accelerated by radiation pressure and to move away, expanding and dispersing over months, as the object gradually returns to its maximum visual brightness. The mechanism of the dust cloud formation and its temporal evolution are, however, still poorly understood. Especially, the location of dust formation is in dispute: far from the star, $\\ga 20 \\mbox{$R_{\\star}$}$ (e.g. Fadeyev \\cite{fadeyev86}, \\cite{fadeyev88}, Feast \\cite{feast96}), or very close to the photosphere, $\\sim 2 \\mbox{$R_{\\star}$}$ (Payne-Gaposchkin \\cite{PG63}). Recent photometric and spectroscopic observations as well as theoretical progress on dust formation suggest that the latter scenario may be the case (see, e.g. Clayton \\cite{clayton96}, Feast \\cite{feast97ii}), however, no definitive answer is yet available.  The spectral energy distributions (SEDs) of RCB stars exhibit infrared emission peaks around 6 $\\sim$ 8~\\mbox{$\\mu$m}. The IR excess, which accounts for typically 30\\% of the total flux, is constantly present, regardless of the visual brightness of the central star. Therefore, the IR excess originates not from a single newly formed dust cloud, but mainly from a group of dispersed dust clouds with temperatures of approximately 600 -- 900~K. For example, Walker et al. (\\cite{walker96}) fit the infrared (spectro)photometric data of R CrB with a 650~K blackbody. The study of IRAS observations by Gillett et al. (\\cite{gillett86}) led to the detection of an additional, very extended ``fossil'' shell around R CrB, whose diameter is as large as 18\\arcmin, with a temperature of $\\sim 30$~K. Walker (\\cite{walker94}) found such fossil shells for at least four RCB stars.   Apart from the detection of the fossil shells, most of the observational results on the circumstellar environment around RCB stars were obtained by photometry and spectroscopy.  Recently, Clayton \\& Ayres (\\cite{clayton01}) have revealed extended \\ion{C}{ii} $\\lambda$1335 emission around two RCB stars, \\object{V854~Cen} and \\object{RY~Sgr}, by long-slit spectroscopy. However, such direct information on the spatial distribution of material in the vicinity of the central star has been very rare up to now. In this paper, we present high-resolution speckle interferometry carried out for R~CrB at maximum and minimum light. The properties of the warm dust shell will be derived by simultaneous fits of the observed visibilities and SEDs using power-law models. We will also discuss the possible indication of a newly formed hot dust cloud. ",
        "conclusions": "Our 75~mas resolution speckle interferometric observations with the SAO 6~m telescope have spatially resolved the dust shell around R CrB for the first time. Neither the observation near maximum light nor at minimum light shows any clear deviation from circular symmetry.  In order to derive the size of the dust shell, we first considered models consisting of the central star and an optically thin dust shell, neglecting the thermal emission of a newly formed dust cloud. Simultaneous fits of the models to the observed SED and visibility have demonstrated that a model with the central star and an optically thin dust shell with density proportional to $r^{-2}$ seems to be appropriate for R CrB near maximum light. The inner boundary is found to be $82~\\mbox{$R_{\\star}$}\\ \\mbox{(19~mas)} \\pm 23$~\\mbox{$R_{\\star}$}\\ with a temperature of $920 \\pm 103$~K.  This simple picture fails to simultaneously reproduce the SED and visibility observed at minimum light, which has led us to investigate models with thermal emission from a newly formed optically thick dust cloud whose angular size is not yet large enough to be spatially resolvable.  The SED and visibility obtained at minimum light were shown to be well fitted with such models. The presence of a newly formed dust cloud as hot as 1200~K with a radius of 4 -- 5~\\mbox{$R_{\\star}$}\\ is inferred, together with an optically thin dust shell with $\\mbox{$r_{\\rm in}$} \\sim 170$~\\mbox{$R_{\\star}$}\\ and $\\mbox{$T_{\\rm in}$}\\ \\sim 690$~K. Furthermore, the SED and visibility obtained near maximum light were shown to be fitted also using a model with a newly formed dust cloud {\\em out of } the line of sight. However, we have also discussed that the discrepancy found for the minimum light data may be attributed to the unusual extinction curve of the obscuring dust cloud. Observations during the very bottom of a deep minimum, when the contribution of the central star is truly negligible, are crucial for investigating the dust shell and dust clouds.   \\begin{acknowledgement} We have used, and acknowledge with thanks, data from the AAVSO International Database, based on observations submitted to the AAVSO by variable star observers worldwide. N.R.I. acknowledges the support of the Long-Term Cooperation program of the Alexander von Humboldt Foundation.    \\end{acknowledgement}"
    },
    "0201/astro-ph0201080_arXiv.txt": {
        "abstract": "The November 1999 outburst of the transient pulsar SAX~J2103.5+4545 was monitored with the large area detectors of the Rossi X-Ray Timing Explorer until the pulsar faded after a year. The 358 s pulsar was spun up for 150 days, at which point the flux dropped quickly by a factor of $\\approx $ 7, the frequency saturated and, as the flux continued to decline, a weak spin-down began. The pulses remained strong during the decay and the spin-up/flux correlation can be fit to the Ghosh \\& Lamb derivations for the spin-up caused by accretion from a thin, pressure-dominated disk, for a distance $\\approx 3.2$ kpc and a surface magnetic field $\\approx 1.2 \\times 10^{13}$ Gauss. During the bright spin-up part of the outburst, the flux was subject to strong orbital modulation, peaking $\\approx 3$ days after periastron of the eccentric 12.68 day orbit, while during the faint part, there was little orbital modulation. The X-ray spectra were typical of accreting pulsars, describable by a cut-off power-law, with an emission line near the 6.4 keV of K$_{\\alpha}$ fluorescence from cool iron. The equivalent width of this emission did not share the orbital modulation, but nearly doubled during the faint phase, despite little change in the column density. The outburst could have been caused by an episode of increased wind from a Be star, such that a small accretion disk is formed during each periastron passage. A change in the wind and disk structure apparently occurred after 5 months such that the accretion rate was no longer modulated or the diffusion time was longer. The distance estimate implies the X-ray luminosity observed was between $1 \\times 10^{36}$ ergs s$^{-1}$ and $6 \\times 10^{34}$ ergs s$^{-1}$, with a small but definite correlation of the intrinsic power-law spectral index. ",
        "introduction": "The transient X-ray source SAX J2103.5+4545 was discovered by the Wide Field Camera instrument on the BeppoSAX X-ray satellite during the outburst between February and September 1997 (Hulleman, in 't Zand, $\\&$ Heise 1998). The source showed 358.61 s  pulsations. The X-ray spectrum was consistent with a power-law model. The photon index was 1.27$\\pm$0.14 and the absorption column density was 3.1$\\pm$1.4 $\\times 10^{22}$ cm$^{-2}$.  Another outburst was detected 2 years later by the All-Sky Monitor (ASM) on the Rossi X-Ray Timing Explorer (RXTE). Pointed observations were carried out with  RXTE. Doppler shifts of the pulsations seen in these observations revealed that the orbital period is 12.68 days (Baykal, Stark $\\&$ Swank 2000a,b). The orbital parameters suggest that the source has a high mass companion, but no suitable optical counterpart has been reported. Hulleman et al. (1998) pointed out a B star at the edge of BeppoSAX  error box, but its distance would imply a luminosity too low to explain the spin-up that was seen in the initial RXTE observations.  SAX J2103.5+4545 continued to be active more than a year after the ASM detectors detected it in 1999 November. During the active interval the source was  monitored through regular pointed RXTE observations. In this work, we present results from analysis of the X-ray spectra as well as new results from the pulse timing of the full set of observations. During this outburst, although the source is not a bright target ( $\\le 30$ mCrab), it was possible to make very significant measurements which have a bearing on several aspects of accreting pulsars. ",
        "conclusions": "A correlation between spin-up rate and X-ray flux in different energy ranges has been observed in outbursts of 5 transient systems. These systems are EXO 2030+375, (Parmar, White $\\&$ Stella 1989, Parmar et al. 1989, Reynolds et al. 1996), 2S 1417-62 (Finger, Wilson $\\&$ Chakrabarty 1996), A 0535+26 (Bildsten et al. 1997, Finger, Wilson $\\&$ Harmon 1996), GRO J1744-28 (Bildsten et al. 1997) and XTE J1543-568 (In 't Zand, Corbet $\\&$ Marshall 2001). All these sources were observed during  spin-up phases and the correlations between spin-up rate and X-ray luminosity were explained in terms of accretion from an accretion disk. The outburst of SAX J2103.5+4545 started with a spin-up trend, made a transition to a steady spin rate  and then appeared to just begin a spin-down trend (see figure 7). These measurements of SAX J2103.5+4545 are the first which have resolved the transition to spin-down.   If the accretion is from a Keplerian disk, at the inner disk edge the magnetosphere disrupts the Keplerian rotation of the disk, forcing matter to accrete along magnetic field lines. The inner disk edge $r_{o}$ moves inward with increasing mass accretion rate. The dependence of the inner disk edge $r_{o}$ on the mass accretion rate $\\dot M$ may be approximately expressed as (Pringle $\\&$ Rees 1972, Lamb, Pethick, $\\&$ Pines 1973) \\begin{equation} r_{o}= K \\mu^{4/7}(GM)^{-1/7}\\dot M^{-2/7} \\end{equation} where $\\mu$=$BR^{3}$ is the neutron star magnetic moment with $B$ the magnetic field and  $R$ the neutron star radius, G is the gravitational constant, and $M$ is the mass of the neutron star. In this equation $K=0.91$ gives the Alfven radius for spherical accretion. Then the torque estimate is given by Ghosh \\& Lamb (1979) as \\begin{equation} 2\\pi I \\dot \\nu = n(w_{s}) \\dot M~l_{K}, \\end{equation} where $I$ is the moment of inertia of the neutron star, $l_{K} = (GM)r_{o})^{1/2}$ is the specific angular momentum added by a Keplerian disk to the neutron star at the inner disk edge $r_{o} $; \\begin{equation} n(w_{s}) \\approx 1.4 (1-w_{s}/w_{c})/(1-w_{s}) \\end{equation} is a dimensionless function that measures the variation of the accretion torque as estimated by the fastness parameter \\begin{equation} w_{s} =\\nu /\\nu _{K}(r_{o}) = (r_{o}/r_{co})^{3/2} = 2 \\pi K^{3/2} P^{-1}  (GM)^{-5/7} \\mu ^{6/7} \\dot M^{-3/7}, \\end{equation} where $r_{co}=(GM/(2\\pi\\nu)^{2})^{1/3}$ is the corotation radius at which the centrifugal forces balances the gravitational forces, $w_{c}$ is the critical fastness parameter at which the accretion torque is expected to vanish. The critical fastness parameter $w_{c}$ has been estimated to be $\\sim$ 0.35 and depends on the electrodynamics of the disk (Ghosh \\& Lamb 1979, Wang 1987, Ghosh 1993, Torkelsson 1998).  The behavior of the dimensionless function $n$ as a function of $\\omega_{s}$ can be understood as follows. The accretion torque is the sum of the torque produced by accretion of the angular momentum of the matter that falls onto the star (mechanical torque) and the torque contributed by the twisted magnetic field lines from the star that interact with the outer parts of the disk (magnetic torque). The mechanical torque always acts to spin-up a star rotating in the same sense as the disk flow, whereas torque from the magnetic stresses can have either sign, since the azimuthal pitch of the stellar magnetic field lines that interact with Keplerian flow in the disk changes sign at the corotation radius $r_{co}$. The torque from the magnetic field lines threading the disk between the inner disk edge $r_{o}$ and corotation radius $r_{co}$ is positive, whereas the contribution of torque from the magnetic field lines threading the disk outside the corotation radius $r_{co}$ is negative. Therefore, either spin-up or spin-down torque is possible in this model as a net effect of the balance between the two contributions. Equation 4 gives an analytic expression approximating numerical calculations of the dimensionless torque (Lamb 1988, Ghosh 1993). The torque will cause a spin-up if the neutron star is rotating slowly ($w_{s}~<~w_{c}$) in the same sense as the circulation in the disk. Even if the neutron star is rotating in the same sense as the disk flow, the torque will be in the direction of spin-down if the neutron star is rotating too rapidly ($w_{s}~>~w_{c}$).  The accreted material will produce X-ray luminosity at the neutron star surface at the rate \\begin{equation} L =  GM \\dot M /R \\end{equation} From equations 2,3 and 6, the rate of spin-up is related to the X-ray intensity through \\begin{equation} \\dot \\nu \\propto n(w_{s}) L^{6/7} =  n(w_{s}) (4 \\pi d^{2} F)^{6/7}, \\end{equation} where $d$ is the distance to the source and $F$ is the X-ray flux.  Since the source distance and magnetic field are not known, we fit the Ghosh \\& Lamb model of the pulse frequency derivative for these two parameters. Because the power index between the spin-up rate and  the X-ray luminosity is a result of the theoretical model, as a test of the model we also fit for this index, obtaining 0.75 $\\pm$ 0.13, which is consistent with the value 6/7 expected in the model. We obtain for the distance to the source  3.2 $\\pm$ 0.8 kpc and for the magnetic field (12 $\\pm$ 3) $\\times 10^{12}$ Gauss. As seen in figure 6, the points corresponding to spin-up and spin-down are consistent with falling on the same curve. With the estimated distance of 3.2 $\\pm$ 0.8 the source could be in a star formation region in the Perseus arm at approximately  4 kpc (Vogt \\& Moffat 1975, Georgelin \\& Georgelin 1976. The X-ray flux at the peak of the outburst $\\sim 6.4 \\times 10^{-10}$ erg s$^{-1}$ cm$^{-2}$ corresponds to an accretion luminosity of $\\sim 8.8 \\times 10^{35}$ erg s$^{-1}$ and a mass accretion rate $ \\sim 6.5 \\times 10^{15}$ g s$^{-1}$). The magnetic field estimate $\\sim 12 \\times 10 ^{12}$ Gauss corresponds to a cyclotron emission line at $\\sim$ 140 keV, which is too high to be detected by HEXTE, because of the low count rate of the source at this energy range.  During the decrease of mass accretion rate $\\dot M$, the fastness parameter increases as $w_{s} \\sim \\dot M^{-3/7}$. When the fastness parameter reaches $w_{s}=(r_{o}/r_{co})^{3/2} \\sim $ 0.35, the inner edge of the disk will be still be well inside the corotation radius $r_{o} \\sim 0.5 r_{co}$. In this case, even if the net torque on the neutron star vanishes, a large pulse fraction change is not expected, since the polar cap region at which the material accretes and radiates will not be changed significantly. Figure 8, presents the 3-20 keV pulse profiles at the bright and faint phases of the outburst. It is clearly seen that there is no significant changes in the pulse profiles.   Further decrease of the mass accretion rate eventually will expand the magnetospheric radius $r_{m} \\approx  r_{o}$ to the corotation radius $r_{m} \\sim r_{co}$. Then, some of the material will be accelerated to super-Keplerian velocities and can not easily be accreted. It may be expelled from the system. Accretion of material will carry angular momentum and tend to spin up the neutron star, while the expulsion of matter will extract angular momentum from the star. These forces tend to bring the neutron star into rotation  at the equilibrium period. It is expected that accretion is eventually centrifugally inhibited. In this propeller regime the neutron star would rapidly spin-down and the X-ray luminosity might be produced by the release of gravitational energy at the magnetosphere (King $\\&$ Cominsky 1994, Campana et al. 1995). X-ray luminosity from the magnetospheric emission would be reduced by a factor of $10^{3}-10^{4}$ (Corbet 1996).  We have assumed here that the accretion is still through a disk. Given that the luminosity is very low ($< 10^{36}$ ergs s$^{-1}$) and the companion probably an early type star, it is not obvious that the accreting material would be able to form a disk, that is, have enough angular momentum to circularize outside the magnetosphere. Wang (1981) showed that this would be the case if the wind velocity is relatively slow ($< 500$ km s$^{-1}$) rather than fast ($> 1000$ km s$^{-1}$). The companion is likely to be a Be star, although not yet identified, because these have episodes of dense slow winds in the equatorial plane (Waters at al. 1988; Li $\\&$ van den Heuvel 1996).  However, if the wind is slow enough that the velocity relative to the neutron star is dominated by the orbital velocity, the accretion would be relatively independent of phase or tend to peak at apastron rather than near periastron as observed. Diffusion through a disk might shift the time of peak X-ray luminosity, but if the shift is as long as a week, the diffusion would also reduce the amplitude of modulation. It is a fast wind with a spherical outflow for which the rate of capture would peak at periastron.  What is observed appears to be the opposite of a simple picture of an episode of equatorial slow wind decaying to a fast wind. The phasing of the modulation shown by SAX J2103.5+4545 has been seen for other sources, such as V0332+53 (Waters et al. 1989), in one outburst in which it appeared consistent with prompt response to a wind of about 300 km s$^{-1}$.  The simple picture does not capture the complexity of Be Star winds. The influence of a wind rotational velocity and a range of velocity laws (or density dependence on distance from the star), in addition to the wind velocity, were studied by Waters et al. (1989). For the SAX~J2103.5+4545 outburst that we monitored, a wind velocity of 200-300 km s$^{-1}$ and a typical mass loss rate of 10$^{-8} \\msun$ yr$^{-1}$ of a Be star could give the implied peak X-ray luminosities, the orbital modulation, and the phase dependence. But the additional parameters would be needed and several different regimes seem possible.    The drop in accretion rate which we observed in SAX~J2103.5+4545 coincided with change in the rate of angular momentum exchange. Presumably, both are caused by changes in the wind structure. Exactly what these are is not obvious. The smooth behavior of the angular acceleration of the pulsar, as opposed to erratic (random walk) variations like those of Vela X-1 (Bildsten 1997), for example, implies the accretion is still via a disk.  The accretion rate has not yet fallen so low that the propellor effect comes into play.  The disk accretion should moderate orbital dependence. For there to be a disk and yet strong orbital modulation at the peak of the outburst, the disk must be small. The lack of modulation at the end of the outburst would be consistent with a disk that is more extensive, through which the diffusion is longer.   Li and van den Heuvel (1996) considered the spin period versus binary period in the diagram first constructed by Corbet (1984). The accretion from a slow wind to a neutron star spinning in equilibrium appears responsible for the main correlation of the Be star systems. Accretion from a fast wind is manifested by supergiants with longer pulse periods for a given orbital period and for Be stars in certain phases. The values of spin and orbital period place SAX~J2103.5+4545 among these sources, but a supergiant is usually a steady rather than a transient source. SAX~J2103.5+4545 can be viewed as not having yet achieved equilibrium. The transient episode decreased the pulse period by $\\approx 0.9$ s. If this occurs every 1.4 yr, the source would spin up to the equilibrium line (1--2 s for a 13 d orbital period) in $< 600$ yr. The frequency value obtained in the 2001 outburst was $1.4 \\times 10^{-3}$ mHz {\\em above} the last previous measurement. It is consistent with an average spin-up rate during the first 30 days (unobserved) of the ourburst being an average of $5 \\times 10^{-13}$ Hz s$^{-1}$, about the same as in the beginning of the 1999 outburst. The spin-down between the last monitored value in 2000 and the beginning of the 2001 outburst would have been only about $10^{-4}$ mHz at the rate observed. The spin rate apparently ratchets up with each outburst.   The X-ray spectra in principal can carry information about the disk and the wind as well as about the flow onto the neutron star. The 6.4 keV line, which comes from fluorescence from relatively cold Fe, could come from all of these sources. Fluorescent Fe K$_{\\alpha}$ is common to many high-magnetic-field  binary pulsars. The equivalent widths in the bright and faint phase spectra of SAX J2103.5+4545 are similar in magnitude to those observed in spectra of Vela~X--1 and GX~301--2 (summarized in Nagase 1989), the accretion rate in SAX J2103.5+4545 is 10--100 times lower and there is no varying column density indicating formation of a shell around the source. Yet for Vela X-1, where the neutron star is eclipsed,  the behavior of the line as a function of binary phase implied a source close to the neutron star in addition to the wind. In the case of SAX~J2103.5+4545, the lack of orbital dependence of the equivalent width during the bright part of the outburst again suggests the line is produced close to the neutron star where the geometry of cool material can be independent of the orbital position. In the decaying part of the outburst, the equivalent width is nearly twice as large. The moderate column density implies the material is seen in reflection rather than transmission. The cool fluorescing material may subtend a  larger solid angle, or a second component is excited for the new type of wind. The data could not significantly limit  orbital modulation of the line emission during this part of the outburst.   The spectral index of the intrinsic X-ray spectrum appears to be consistently correlated with the flux, harder for higher flux, both in the variation with orbital phase and the variation during the outburst. The fact that the pulse amplitude and shape did not change suggests that the X-rays continue to be produced in flow to the neutron star and in that case the X-ray flux seems likely to be proportional to the rate of accretion onto the surface.  Intrinsic X-ray spectra for low luminosity sources $\\dot M < 10^{17}$ g s$^{-1}$ have been modelled (Meszaros et al. (1983); Harding et al. (1984)). The continuum spectrum has not appeared to be very sensitive to the mass accretion rate, but the contributions of the cyclotron line emission were not included and could make a difference if Compton scattering of line photons affects the continuum below the line energy. There has appeared to be at least a relation between the cut-offs observed in pulsar spectra and the cyclotron resonance energies (Makashima \\& Mihara 1992; see however Mihara, Makashima, \\& Nagase 1998). The high magnetic field obtained from the spin rate dependence puts the cyclotron resonance energy farther above the cut-off energy than is the case for pulsars with identified resonance features and would give SAX~J2103.5+4545 the highest field of the accreting pulsars (Orlandini \\& Dal Fiume 2001). The magnetic field estimate depends on the critical fastness and the balance between the magnetosphere and the accretion flow. Quantitative estimates seem likely to be subject to dependence of the model on assumptions about the nature of the disk, for example. Understanding the X-ray spectra and flux correlations appears to require better models of the spectral formation and also the magnetospheric physics.  RXTE's monitoring of this transient pulsar allowed the spin rate to be tracked while the luminosity declined to $6 \\times 10^{34}$ ergs s$^{-1}$. The spin-rate responded smoothly to the flux, spinning up, coming to equilibrium and then reversing sign, just as the Ghosh \\& Lamb characterization of the magnetospheric interaction predicted. While the magnetic field and distance that come out of fitting the data may be estimates subject to model dependence of the spin-up/flux relation, this source demonstrates the possibility of determining system parameters from such observations. The system is also interesting in that it appears to be  different from better known ones and provides information about uncharted parts of parameter space of pulsars in binaries. Detailed study of the orbital phase dependence during the outbursts of this source will provide constraints on the stellar orbit flows from which this neutron star accretes."
    },
    "0201/astro-ph0201159_arXiv.txt": {
        "abstract": "{We present deep $BVI$ observations of the dwarf irregular galaxy UKS1927-177 in Sagittarius (SagDIG). Statistically cleaned $V$, $(B-I)$ color-magnitude diagrams clearly display the key evolutionary features in this galaxy. Previously detected C stars are located in the color-magnitude diagrams and shown to be variable, thus confirming the presence of a significant upper-AGB intermediate age population. A group of likely red supergiants is also identified, whose magnitude and color is consistent with a 30 Myr old burst of star formation. The observed colors of both blue and red stars in SagDIG are best explained by introducing a differential reddening scenario in which internal dust extinction affects the star forming regions. Adopting a low reddening for the red giants, $E(B-V) = 0.07 \\pm 0.02$, gives [Fe/H]=$-2.1 \\pm 0.2$ for the mean stellar metallicity, a value consistent with the [O/H] abundance measured in the \\hii\\ regions. This revised metallicity, which is in accord with the trend of metallicity against luminosity for dwarf irregular galaxies, is indicative of a ``normal'', although metal-poor, \\di\\ galaxy. A quantitative description is given of the spatial distribution of stars in different age intervals, in comparison with the distribution of the neutral hydrogen.  We find that the youngest stars are located near the major peaks of emission on the \\hi\\ shell, whereas the red giants and intermediate-age C stars define an extended halo or disk with scale length comparable to the size of the hydrogen cloud.  The relationship between the distribution of \\ism\\ and star formation is briefly discussed. ",
        "introduction": "\\label{s_intro}  The Sagittarius dwarf irregular (\\abbrev{\\di}) galaxy, also known as SagDIG or UKS~1927-177, is a quite difficult object to study because of its low Galactic latitude and consequent high foreground contamination.  First reported by Cesarsky \\etal (\\cite{cesarsky77}) on ESO Schmidt photographic plates, the galaxy was studied by Longmore \\etal (\\cite{longmore78}), who derived a total luminosity $M_B=-10.5$. The first CCD study of the resolved star content in SagDig was that of Cook (\\cite{cook87}; hereafter C87), who obtained both intermediate-band and broad-band photometry. Recently, two new investigations have improved our knowledge of the distance and stellar content of this galaxy (Karachentsev \\etal \\cite{kara99}; Lee \\& Kim \\cite{lee00}; hereafter KAM99 and LK00 respectively). A new distance modulus was derived from the tip of the \\rgb, equal to $(m-M)_0=25.13$ in KAM99 and $(m-M)_0=25.36$ in LK00, respectively. This distance indicates that SagDIG is a member of the Local Group (LG), confirming the evidence from its negative radial velocity and position in the $V_{\\odot}$ $vs.$ $\\cos{\\theta}$ diagram (van den Bergh \\cite{vandenberg94}; Pritchet \\& van den Bergh \\cite{pritchet99}). Both studies, while adopting different estimates for the reddening, concluded that SagDIG has the lowest metallicity among the LG star forming dwarf galaxies (KAM99 estimated [Fe/H]$\\sim-2.45$, while LK00 give [Fe/H] in the range $-2.8$ to $-2.4$; both estimates are based on an extrapolation of the calibration relation for Galactic globular clusters).  The cold and warm interstellar medium (\\abbrev{\\ism}) of SagDIG and its kinematics have been the subject of several investigations. Skillman \\etal (\\cite{stm89}) obtained optical spectrophotometry of the most luminous \\hii\\ regions, and estimated a [O/H] abundance $\\sim3\\%$ of the solar value. New measurements of the O and N abundance in the \\hii\\ regions of SagDIG are presented in a companion paper (Saviane et al. \\cite{savi+01}).  The new estimate, $12+\\mbox{log(O/H)}=7.23\\pm 0.20$, is by 0.2 dex more metal-poor than found by Skillman et al. (\\cite{stm89}). The photometric properties of the \\hii\\ regions were investigated by Strobel \\etal (\\cite{strobel91}), who detected 3 regions in their $2\\farcm5$ square field.  High-resolution, high sensitivity VLA observations of the SagDIG \\hi\\ content have been obtained by Young \\& Lo (\\cite{young97}) (see also Lo et al. \\cite{lo+sar}).  About $1.2 \\times 10^7$ \\msol\\ of H$+$He have been estimated (using the new distance from KAM99 and the \\hi\\ mass from Young \\& Lo), distributed in an almost symmetric ring likely produced by the combined effects of stellar winds and supernovae. SagDIG thus appears to have a high mass fraction in the form of neutral gas, $M_{H I}/L_B= 1.6$, a value quite typical of the \\hi\\ content in dwarf irregular galaxies (see Mateo \\cite{mateo98}).  Because of its high gas content, low luminosity, and especially its claimed very low metallicity, SagDIG may be a clue to the origin and evolution of dwarf galaxies. Although located at the border of the Local Group it is still close enough to allow us to study both its stellar populations and the gaseous component, and to compare their physical properties. In particular, a sound knowledge of the metallicity of SagDIG is important to constrain the luminosity-metallicity relation for {\\di}s at the metal-poor end, and to trace the metal enrichment history in dwarf galaxies. All this motivated an independent study of the metallicity, distance, and stellar content of SagDIG. The large baseline provided by the $(B-I)$ color, together with an analysis based on statistical subtraction of the Galactic foreground, allowed us to improve the discrimination of young and old stellar populations in the color-magnitude diagram (CMD). By accounting for the differential effects of internal dust extinction, we revise upward the [Fe/H] estimate and show that the metal abundances of the red stars and the \\ism\\ can be easily reconciled.  The plan of the paper is as follows. In Sect.~\\ref{s_observ} we present the data reduction and calibration. Section~\\ref{s_cmd} presents color-magnitude diagram in different colors, using a statistical correction of the foreground contamination. The CMD location of C stars is also discussed. In Sect.~\\ref{s_basic} we rederive the distance and metallicity of SagDIG by assuming a different reddening for the young and old populations. The properties and different spatial distributions of the blue and red stellar populations are quantified and discussed in Sect.~\\ref{s_stelpop}, where the surface density of young stars in different age intervals is compared with the projected distribution of neutral hydrogen. As a result of the revised metallicity, SagDIG is shown to fit well the known luminosity-abundance trend of dwarf irregular galaxies. Our results are summarized in Sect.~\\ref{s_summary}. ",
        "conclusions": "\\label{s_summary}  We have presented a $BVI$ study of the stellar content and metallicity of UKS~1927-177, the Sagittarius dwarf irregular galaxy, based on deep color-magnitude diagrams.  We have found that {\\it the color-magnitude diagram of SagDIG is better understood by introducing a differential reddening scenario} in which the young stars, located near high-density \\hi\\ clumps, are more reddened than the older stars distributed all over the galaxy.  In fact, the color distribution of the Galactic foreground stars indicates a relatively low reddening (we assume $E_{B-V}=0.07 \\pm 0.02$ for the old stars), while measurements of the Balmer decrement in \\hii\\ regions (Skillman et al. \\cite{stm89}; Saviane et al. \\cite{savi+01}) provide a higher reddening, $E_{B-V}=0.19$. Using the lower reddening for the ``old'' population in SagDIG, we obtained revised values for the metallicity and distance, [Fe/H]$=-2.1 \\pm 0.2$ and $(m-M)_{0}=25.14 \\pm 0.18$, respectively.  While the distance confirms previous estimates, {\\it our metallicity turns out to be significantly higher} than those proposed by Karachentsev et al. (\\cite{kara99}) and Lee \\& Kim (\\cite{lee00}). {\\it Using this new metallicity, we have compared SagDIG with other dwarf galaxies in the luminosity-metallicity diagram, and found that it is consistent with the general trend for \\di\\ galaxies.}  The large baseline given by the $(B-I)$ color, together with an analysis based on statistical subtraction of the Galactic foreground, provided additional information on the recent star formation in SagDIG.  The C stars from two sources have also been compared with our cleaned CMDs, which confirmed the presence of an intermediate age population. For the young stellar population, we have adopted the reddening derived for the \\hii\\ regions, thus obtaining a good match to the theoretical isochrones. A comparison with models indicates a significant burst taking place $\\sim$30 Myr ago, with star formation going on until $\\sim$10 Myr ago. This burst seems to have interrupted a relatively quiescent period between $\\sim$30--100 Myr ago. We identified a group of candidate red supergiants that are quite well fitted by the isochrones of young He-burning stars.  We have also obtained a detailed quantitative comparison of {\\it the spatial distribution of stars in different age ranges} with the distribution of the ISM. Different {density profile scale lengths} have been measured for the young blue stars and the red giants. The youngest stars have been shown to concentrate in a central region nearly coincident with the density peaks in the \\hi\\ distribution. In constrast, the distribution of red giants is quite extended, yet it seems to be correlated with the HI morphology as well. This may suggest for the red giants in SagDIG a contribution by relatively young (a few Gyr old) stars. We also noted that some blue stars are found in ``tails'', perhaps a hint of spiral structures, quite far from the main star-forming sites. This suggests a similarity with the small spiral-like structures that can form even in a small galaxy according to the theory of stochastic self-propagating star formation (Gerola \\& Seiden \\cite{gerola78}).  An alternative possibility is that these features are a vestige of star formation related to the gaseous shell."
    },
    "0201/astro-ph0201473_arXiv.txt": {
        "abstract": "{\\sl RXTE\\/} observations confirm that the X-ray lightcurve of \\voph\\ is pulsed at the beat cycle, as expected in a discless intermediate polar. There are no X-ray modulations at the orbital or spin cycles, but optical line profiles vary with all three cycles.  We construct a model for line-profile variations in a discless accretor, based on the idea that the accretion stream flips from one magnetic pole to the other, and show that this accounts for the observed behaviour over the spin and beat cycles. The minimal variability over the orbital cycle implies that 1) \\voph\\ is at an inclination of only \\appro 10\\deg, and 2) much of the accretion flow is not in a coherent stream, but is circling the white dwarf, possibly as a ring of denser, diamagnetic blobs. We discuss the light this sheds on disc formation in intermediate polars. ",
        "introduction": "The magnetic cataclysmic variables can be divided into two classes. In polars (or AM~Her stars) the magnetic field of the white dwarf is strong enough to lock its rotation to the binary orbit, and accretion proceeds via a stream which is deflected by the field onto a magnetic pole. In contrast, in intermediate polars (IPs or DQ~Her stars) a lower-field white dwarf spins more rapidly than the orbit, and the stream feeds into an accretion disc, which then feeds field lines from its inner edge (see Warner 1995 for a comprehensive review).  The distinction is blurred, however, by the existence of several asynchronous polars, in which the spin periods differ from the orbit by \\sqig 1 per cent (e.g.\\ Schwope \\etal\\ 1997). There has also been a long debate on the existence of discless IPs (e.g.\\ Hameury, King \\&\\ Lasota 1986; King \\& Lasota 1991; Hellier 1991; Wynn \\&\\ King 1992).  In such systems, the stream would be expected to flip first to the upper pole, then to the lower pole, as the magnetic dipole rotates (e.g.\\ Hellier 1991; Wynn \\&\\ King 1992). Such pole flipping would occur at the frequency with which the relative geometry changes, namely the beat frequency \\beat, where $\\Omega$ and $\\omega$ are the orbital and spin frequencies respectively.  Buckley \\etal\\ (1995; 1997) reported the first secure evidence for a discless IP, with the discovery of the {\\it Rosat\\/} source \\voph\\ (RX\\,J1712.6--2414). They found that polarised light from the system varies at 927 s, which is interpreted as the spin period of the magnetic dipole and thus of the white dwarf.  The X-rays, though, are pulsed at 1003 s, which is the beat period between the 3.42-h orbital and 927-s spin cycles; thus \\voph\\ shows the signature of pole-flipping accretion, as expected in a discless IP.  \\voph\\ is a prime opportunity to study the interaction of an accretion stream with a magnetic field in a situation where the flow is continually changing as the dipole rotates. This contrasts with most magnetic cataclysmic variables where the flow is expected to settle (at least temporarily) into a quasi-equilibrium.  (In principle the asynchronous polars offer the same opportunity, but it is much harder to obtain good coverage of their \\sqig 50-d beat cycles.)  In this paper we present new X-ray observations of \\voph, along with optical spectroscopy aimed at tracing the accretion flow as it connects to the field lines.   \\begin{figure*}\\vspace*{5.5cm}     % \\caption{The 2--15-\\kev\\ X-ray lightcurve of \\voph\\ as recorded by {\\sl RXTE}.} \\special{psfile=v24f1.eps hoffset=-75 voffset=-152 hscale=80 vscale=80 angle=0} \\end{figure*}  \\begin{figure*}\\vspace*{4.8cm}     % \\caption{The Fourier transform of the {\\sl RXTE\\/} observation, part of which is shown in Fig.~1. The orbital ($\\Omega$) and spin ($\\omega$) frequencies are marked. The dominant modulation, with an alias structure caused by the spacecraft orbit, is at the beat frequency, \\beat.} \\special{psfile=v24f2.eps hoffset=-60 voffset=-130 hscale=75 vscale=75 angle=0} \\end{figure*}  \\voph\\ is the most polarised IP, leading to a magnetic field estimate of 9--27 MG (Buckley \\etal\\ 1995; V\\\"ath 1997), the highest for any IP and overlapping with the range for low-field polars. The high field may explain why a disc has not formed in this system, whereas they do in IPs at \\sqig 1 MG.  Buckley \\etal\\ (1995) also found that the circular polarisation is always of the same sign, and so concluded that we only ever see one magnetic pole. Given, also, a lack of radial velocity motion at the orbital period, they proposed that \\voph\\ is at a low inclination, and that we only see the pole nearest us (the `upper' pole). From the fact that there were no detectable variations in linear polarisation, Buckley \\etal\\ (1995) suggested that this pole was never viewed side on, and so proposed that both the inclination of the binary, and the angle between the spin and magnetic axes of the white dwarf, $\\delta$, are low.  \\begin{figure}\\vspace*{7.0cm}     % \\caption{The X-ray lightcurve of \\voph\\ folded on the 1003-s beat period. A typical error is shown.} \\special{psfile=v24f3.eps hoffset=-72 voffset=-98 hscale=57 vscale=57 angle=0} \\end{figure}  \\begin{figure}\\vspace*{4.2cm}     % \\special{psfile=v24f4.ps hoffset=-48 voffset=-100 hscale=54 vscale=54 angle=0} \\caption{An illustration of the discless accretion geometry in \\voph. In our model (Section 4) the stream attaches to a 20\\deg\\ swathe of field lines (shown in bold). The change in geometry over the beat cycle can be visualised by imagining the feeding point moving round the ring.} \\end{figure} ",
        "conclusions": "$\\bullet$ \\voph\\ is at a low inclination, estimated as 10\\deg. We only ever see the upper magnetic pole.  $\\bullet$ There is no accretion disc in this system. The stream-fed accretion flips from pole to pole on the beat cycle, and this can be seen directly in the line profiles. However, only \\appro 25 per cent of the accreting material participates in this motion.  $\\bullet$ The rest of the material appears to be circling the white dwarf, perhaps in the form of diamagnetic blobs.  Such blobs may have been a precursor to disc formation in other intermediate polars.  $\\bullet$ The spin-cycle variations of the emission lines are accounted for by a simple model of stream-fed accretion. The best match between model and data is found for ditances 5--10 \\rwd\\ from the white dwarf. The offset between the magnetic and spin axes of the white dwarf must be at least 30\\deg."
    },
    "0201/astro-ph0201190_arXiv.txt": {
        "abstract": "While the abundances of Be and B observed in metal-poor halo stars are well explained as resulting from spallation of CNO-enriched cosmic rays (CRs) accelerated by supernova shocks, accounting for the observed $^6$Li in such stars with supernova CRs is more problematic. Here we propose that gravitational shocks induced by infalling and merging sub-Galactic clumps during hierarchical structure formation of the Galaxy should dissipate enough energy at early epochs, and CRs accelerated by such shocks can provide a natural explanation of the observed $^6$Li. In clear constrast to supernovae, structure formation shocks do not eject freshly synthesized CNO nor Fe, so that the only effective production channel at low metallicity is $\\alpha-\\alpha$ fusion, capable of generating sufficient $^6$Li with no accompanying Be or B and no direct correspondence with Fe. Correlations between the $^6$Li abundance and the kinematic properties of the halo stars may also be expected in this scenario. Further, more extensive observations of $^6$Li in metal-poor halo stars, e.g. by the Subaru HDS or VLT/UVES, may offer us an invaluable fossil record of dissipative dynamical processes which occurred during the formation of our Galaxy. ",
        "introduction": "The light elements Li, Be and B are unique in that apart from $^7$Li, none can be synthesized appreciably in thermal environments such as stellar interiors, supernova envelopes or the standard Big Bang. Instead, the bulk of these elements are believed to arise from nonthermal nuclear reactions induced by cosmic rays (CRs). Their abundances observed in the Galactic disk, particularly for the isotopes $^6$Li, $^9$Be, and $^{10}$B, are well explained as being products of spallation processes in which CNO atoms in the interstellar medium (ISM) are broken up into LiBeB by collisions with CR protons or $\\alpha$ particles (Reeves, Fowler \\& Hoyle 1970, Meguzzi, Audouze \\& Reeves 1971, Walker, Mathews \\& Viola 1985). In the last decade, extensive observations of LiBeB in population II, metal poor halo stars (MPHS) have turned up new and unexpected results, spurring controversy as to what type of CR sources and production mechanisms were operating in the halo of the early, forming Galaxy (see review by Vangioni-Flam, Cass\\'e \\& Audouze 2000).  To date, most models of light element evolution in the early Galaxy have focused on strong shocks driven by supernovae (SNe) as the principal sources of CRs. Although the assumed CR composition, energy spectrum and the manner of LiBeB production vary in different models, a general consensus is that Be and B in MPHS mainly originate from the ``inverse'' spallation process, whereby CR CNO particles are transformed in flight into LiBeB by impinging on ISM H or He atoms (Duncan, Lambert \\& Lemke 1992). This can be realized if a sizable fraction of the CRs responsible for spallation comprise fresh, CNO-rich SN ejecta (e.g. Cass\\'e, Lehoucq \\& Vangioni-Flam 1995, Vangioni-Flam et al. 2000, Ramaty et al. 2000, Parizot \\& Drury 1999, Suzuki, Yoshii \\& Kajino 1999, hereafter SYK, Suzuki \\& Yoshii 2001, hereafter SY), as opposed to CRs injected from the average ISM (e.g. Fields \\& Olive 1999).  The origin of $^6$Li in MPHS, which has been detected in only 3 stars so far (Hobbs 2000, Nissen 2000 and references therein), is more mysterious, as current models involving SN CRs face some difficulties. A peculiar aspect of Li is that in addition to spallation, the fusion process of CR $\\alpha$ particles with ambient He atoms can be effective, and should actually dominate Li production at low metallicities. (Note that while both $^7$Li and $^6$Li are synthesized in comparable amounts, the CR-produced $^7$Li component is generally overwhelmed by the ``Spite plateau'' from primordial nucleosynthesis in the metallicity range under consideration; e.g. Ryan et al. 2001.) If the CR energy spectrum is taken to be a standard power-law distribution in momentum (\\S 3), one requires a CR injection efficiency much higher than normally inferred to reproduce the $^6$Li observations, whether the CR composition is metal-enriched or not (Ramaty et al. 2000; SY, see their fig.2). Faring better are some models which envision CR acceleration by multiple shocks inside superbubbles, where an additional, $\\alpha$-enriched low energy CR component with a hard spectrum and an energy cutoff at few 100 MeV is present (e.g. Vangioni-Flam et al. 1999, Parizot \\& Drury 1999). However, such a CR component is at the moment only hypothetical, without any observational support. This raises the question of whether there may have been other sources for $^6$Li.  This work investigates a new and more natural $^6$Li production scenario based on a previously unconsidered CR source: CRs accelerated at structure formation shocks, i.e. gravitational virialization shocks driven by the infall and merging of sub-Galactic gas clumps during the hierarchical build-up of structure in the early Galaxy. Such shocks are inevitable consequences in the currently standard theory of hierarchical structure formation. We show below that this picture gives a better explanation of the present data, and also provides a number of testable predictions for future observations that are quite distinct from SN CR models. Our scenario also embodies unique and important implications for understanding the formation of our Galaxy. ",
        "conclusions": "For our calculations, we have selected the following sets of parameters for $t_{\\rm SF}$, $\\tau_{\\rm SF}$ and $\\gamma_{\\rm SF}$, respectively, labeled models I - V: I (0.12, 0.1, 3), II (0.22, 0.1, 3), III (0.32, 0.1, 3), IV (0.22, 0.1, 2) and V (0.1, 0.5, 3), where $t_{\\rm SF}$ and $\\tau_{\\rm SF}$ are in units of Gyr; model VI is the case of SN CRs only. These were chosen to provide results exemplary of light element production by SF CRs, in contradistinction to that by SN CRs. The evolution of $^6$Li and Be vs. metallicity calculated for each model until the end of halo chemical evolution ([Fe/H]$\\simeq -1.5$) is shown in Fig.\\ref{fig:libe}, along with the current observational data for $^6$Li (Smith et al. 1998, Cayrel et al. 1999, Nissen et al. 2000) and Be (Boesgaard et al. 1999) in MPHS.  We discuss some salient points regarding these results. First, it is confirmed that with our fiducial parameters, production by SN CRs alone (VI) works very well for the observed Be (and B, not shown), yet falls short of the observed $^6$Li. Accounting for this by SN CRs demands a much larger value of $E_{\\rm SN}$, $\\xi_{\\rm SN}$ and/or $\\gamma_{\\rm SN}$, which must be supplemented with a larger $f_{\\rm CR}$ in order for Be and B to be consistently reproduced together (see fig.2 in SY).  In contrast, with reasonable values for $\\epsilon_{\\rm SF}$, $\\xi_{\\rm SF}$ and $\\gamma_{\\rm SF}$, production by SF CRs is capable of explaining the current $^6$Li data quite adequately. This mainly owes to two facts: 1) SF shocks are more energetic than (or at least as energetic as) SN shocks, as estimated in \\S 2, and 2) SF CRs can generate $^6$Li at early epochs independently of the metallicity. Regardless of the early evolutionary behavior, identical $^6$Li abundances are attained at the end of the halo phase for a given $\\gamma_{\\rm SF}$ (I - III, V), since this is determined by the time-integrated CR flux, for which we had assumed a fixed value. Compared to a flat spectral index of $\\gamma_{\\rm SF}=2$ (IV), a steeper one of $\\gamma_{\\rm SF}=3$, more appropriate for low Mach number SF shocks (\\S 3), results in a larger $^6$Li yield, by about factor of 3. This is because with a constant total CR energy, a steeper index implies a larger CR flux in the subrelativistic energy range $E \\sim 10$ -- $100$ MeV, where the $^6$Li production cross section peaks (see fig.1 in SY). Note that a conservatively lower specific energy for SF shocks, $\\epsilon_{\\rm SF} \\sim 0.15$ keV/particle (i.e. comparable to SNe), can still be consistent with the available $^6$Li data provided that $\\gamma_{\\rm SF} \\simeq 3$.  As already stressed, 2) is a consequence of $^6$Li synthesis being dominated by $\\alpha-\\alpha$ fusion, and SF shocks not creating any new Fe. Depending on the onset time and duration of the SF shock, the $^6$Li abundance may potentially reach large values quickly at very low metallicity, which can be followed by a plateau or a very slow rise. On the other hand, SNe unavoidably give forth to freshly synthesized Fe, so a correlation with Fe/H must arise; in fact $^6$Li/H vs. Fe/H for SN CRs can never be much flatter than linear (e.g. Vangioni-Flam et al. 1999, Parizot \\& Drury 1999, Fields \\& Olive 1999, Ramaty et al. 2000). Moreover, since SF shocks do not eject fresh CNO, they produce very little Be or B through spallation; only a minuscule contribution can appear as the ISM becomes metal-enriched. This may allow an extremely large $^6$Li/Be ratio at low [Fe/H]. Conversely, we see that SN CRs must play an indispensable role in generating the Be and B observed in MPHS.  Independent of the particular evolutionary parameters, the following abundance trends are characteristic of SF CR production and should serve as distinguishing properties of the scenario for future observations. Going from high to low metallicity: a plateau or a very slow decrease in $\\log ^6$Li/H vs. [Fe/H], followed by a steeper decline in some range of [Fe/H] corresponding to the main epoch of SF; a steady increase in $^6$Li/Be, possibly up to values exceeding $\\simeq 100$, also followed by a downturn. These traits are very distinctive and not expected in SN CR models, for which the slope of $\\log ^6$Li/H - [Fe/H] must be $\\simeq 1$ or greater, and the $^6$Li/Be ratio constant at sufficiently low [Fe/H]. Distinction from any production processes in the early universe (e.g. Jedamzik 2000) should also be straightforward, as they predict a true plateau down to the lowest [Fe/H], in contrast to an eventual decrease for SF CR models. Further, unique diagnostic features are discussed in \\S 5.  Observing $^6$Li in MPHS has proven to be difficult in the past, as measurement of its weak isotopic shift feature relative to the much stronger $^7$Li line requires spectroscopy with very high resolution and signal to noise. The present database of only 3 positive detections in a narrow range of [Fe/H] ($\\sim$ -2.5 -- -2) plus a number of upper limits (Hobbs 2000, Nissen 2000 and references therein) is obviously insufficient for distinction between our SF CR model and various other models. However, the new generation of large aperture telescopes equipped with high resolution spectrographs, e.g. the Subaru HDS or VLT/UVES, should in the near future bring about a larger sample of MPHS with accurately determined $^6$Li/H and $^6$Li/Be over a wide range of [Fe/H], and greatly help us toward deciphering the true origin of $^6$Li in MPHS.     We have put forth a new scenario for the currently puzzling origin of $^6$Li observed in metal-poor halo stars: production by cosmic rays accelerated at structure formation shocks, driven by the hierarchical infall and merging of sub-Galactic structure during the formation of our Galaxy. Several predictions are quite distinct from models involving supernova cosmic rays and clearly testable in the near future, such as the behavior of $^6$Li/H vs. [Fe/H] and $^6$Li/Be vs. [Fe/H] at low [Fe/H], as well as possible correlations between $^6$Li/H and the kinematic properties of halo stars. Since $^6$Li can be construed as a fossil record of dissipative processes during Galaxy formation, further observations of this isotope in halo stars may offer us a unique and invaluable insight into the past dynamical history of our Galaxy."
    },
    "0201/astro-ph0201535_arXiv.txt": {
        "abstract": "The primary X-ray diagnostic lines in He-like ions are mainly excited by electron impact from the ground level to the $n = 2$ levels, but at high temperatures $n > 2$ levels are also excited. In order to describe the atomic processes more completely collision strengths are computed for O~VII including for the first time all of the following: (i) relativistic fine structure, (ii) levels up to the $n=4$, and (iii) radiation damping of autoionizing resonances. The calculations are carried out using the Breit-Pauli R-matrix (BPRM) method with a 31-level eigenfunction expansion. Resonance structures in collision strengths are delineated in detail up to the $n = 4$ thresholds. For highly charged He-like ions radiation damping of autoionizing resonances is known to be significant. We investigate this effect in detail and find that while resonances are discernibly damped radiatively as the series limit $n \\rightarrow \\infty$ is approached from below, the overall effect on effective cross sections and rate coefficients is found to be very small. Collision strengths for the principal lines important in X-ray plasma diagnostics, w,x,y and z, corresponding to the 4 transitions to the ground level $1s^2 \\ (^1S_0) \\longleftarrow 1s2p (^1P^o_1), 1s2p (^3P^o_2), 1s2p (^3P^o_1), 1s2s (^3S_1)$, are explicitly shown. It is found that the effective collision strength of the forbidden z-line is up to a factor of 4 higher at T $< 10^6$ K than previous values. This is likely to be of considerable importance in the diagnostics of photoionized astrophysical plasmas. Significant differences are also found with previous works for several other transitions. This work is carried out as part of the Iron Project-RmaX Network. ",
        "introduction": "Helium-like ions provide the most important X-ray spectral diagnostics in high temperature fusion and astrophysical plasmas. The new generation of X-Ray satellites such as the Chandra X-Ray Observatory and the X-Ray Multi-Mirror Mission-Newton provide high resolution spectra of different types of astronomical objects (e.g. Kaastra \\etal 2000, Porquet and Dubau 2000, Porquet \\etal 2001). The high sensitivity of these observatories and the high quality of the spectra they produce requires highly accurate atomic data for a precise interpretation. The aim of  the Iron Project-R-matrix calculations for X-ray spectroscopy (IP-RmaX) is to calculate extended sets of accurate collision strengths and rate coefficients for all ions of importance in X-Ray diagnostics. Among previous works, the electron impact excitation of Helium like oxygen was previously considered by Pradhan \\etal(1981a,b) in the distorted wave and close coupling approximations for transitions up to the $n=2$ levels. Sampson \\etal (1983) and Zhang and Sampson (1987) used the Coulomb-Born approximation with exchange, intermediate coupling, and some resonances effects to obtain collision strengths for Helium-like ions, with atomic number Z spanning a large range of values ($4<Z<74$). Kingston and Tayal (1983a,b) calculated the collision strength for two transitions, from the ground state to $2 ^3S_1$ and to $2 ^3P^o_1$, using the close coupling R-matrix (RM) method, and derived the corresponding effective collision strengths. Both, the Pradhan \\etal and Kingston and Tayal calculations were in LS coupling. The present work aims at generating a more complete dataset of high reliability for O VII, including all important effects for highly charged ions such as relativistic effects, radiation damping, and resonances in higher complexes up to $n=4$.  The method and computations are summerized in section 2. Results for the collision strengths and important issues are discussed in section 3, and the present results for the effective (Maxwellian averaged) collision strengths are compared with previous calculations. The main conclusions are given in section 4, together with an estimate of accuracy of the results. ",
        "conclusions": "Some of the general conclusions of the paper are as follows.  1. The most complete close coupling calculation using the Breit-Pauli R-Matrix method has been carried out for helium-like oxygen, including resonances up to $n = 4$ levels. Detailed studies of radiation damping indicate that it may have a significant effect on the detailed collision strengths in a small energy region below the threshold(s) of convergence, but not on the effective collision strengths. However, radiation damping is important for higher-Z elements since the transition probabilities increase with Z (Pradhan 1983a,b has found the effect on the z-transition to be 9\\% in $\\Upsilon$ for Fe~XXV at the temperature of maximum abundance of helium-like iron).   2. It is verified that the effects of coupling and resonances from the $n = N+1$ complex play an important role in effective collision strengths for the transitions to the complex $n = N$.  3. The new results for the important z-line transition should significantly affect the analysis of O~VII X-ray spectra from photoionized sources (e.g. active galactic nuclei), where O~VII may be abundant at relatively low temperatures. In collisional ionized (coronal) sources the new results may not affect the theoretically computed line intensities significantly at temperatures close to maximum abundance, but should still do so at lower temperatures. It would be preferable to employ the present data in future collisional-radiative and photoionization models.  4. As all relevant atomic effects in electron-ion collisions  have been considered, and resonances have been carefully delineated, we should expect the present results to be of definitive accuracy. Nonetheless, we conservatively estimate the precision to be about 10-20\\%.  5. All data will be electronically available from the first author from delahaye@astronomy.ohio-state.edu.  The authors would like to thank Dr. Werner Eissner for immense help with many aspects of this work. We also thank Guo-Xin Chen for the new and efficient version of STGF including radiation damping used in these calculations. This work was supported partially by the U.S. National Science Foundation and by NASA. The computational work was carried out on the massively parallel Cray T3E and the vector processor Cray T94 at the Ohio Supercomputer Center in Columbus, Ohio."
    },
    "0201/astro-ph0201229_arXiv.txt": {
        "abstract": "A modification to the Friedmann Robertson Walker equation is proposed in which the universe is flat, matter dominated, and accelerating.  An additional term, which contains only matter or radiation (no vacuum contribution), becomes the dominant driver of expansion at a late epoch of the universe. During the epoch when the new term dominates, the universe accelerates; we call this period of acceleration the Cardassian era.  The universe can be flat and yet consist of only matter and radiation, and still be compatible with observations.  The energy density required to close the universe is much smaller than in a standard cosmology, so that matter can be sufficient to provide a flat geometry.  The new term required may arise, e.g., as a consequence of our observable universe living as a 3-dimensional brane in a higher dimensional universe.  The Cardassian model survives several observational tests, including the cosmic background radiation, the age of the universe, the cluster baryon fraction, and structure formation. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0201/astro-ph0201065_arXiv.txt": {
        "abstract": "More than one third of all nearby galaxies show indications of low-luminosity nuclear activity.  These low-luminosity AGNs (LLAGNs) are traditionally classified into a variety of categories depending on their forbidden-line ratios and on the presence or absence of broad-line emission.  By analogy with Seyfert unification models, it is natural to ask whether the various LLAGN types are different manifestations of the same underlying phenomenon, with observed differences being solely due to orientation and obscuration, or whether there are fundamentally distinct physical processes at work in different categories of LLAGNs.  This contribution reviews some recent observations of LLAGNs in the context of AGN unification scenarios. ",
        "introduction": "During the past two decades, unification models have been remarkably successful at clarifying the underlying connections between a variety of types of AGNs, including Type 1 and 2 Seyferts, radio galaxies and quasars (Antonucci 1993).  These unification scenarios attempt to explain the differences between certain AGN classes as the result of varying orientation to our line of sight of the central engine and an ``obscuring torus'' that surrounds it.  This review explores the question of whether similar unification ideas can also be applied to low-luminosity AGNs (LLAGNs), which make up the vast majority of the AGN population.  These include low-luminosity Seyferts, low-ionization nuclear emission-line regions (LINERs), and transition-type objects with properties intermediate between those of LINERs and those of normal H~II nuclei.  Unification studies of Seyfert galaxies are primarily concerned with figuring out whether Seyfert 2s are intrinsically the same kind of AGNs as Seyfert 1s.  For LINERs, the unification question takes on a somewhat different focus, because it is not clear whether all of the Type 2 LINERs and transition objects are genuine AGNs at all. Ultimately, we may be able to achieve a partial unification in which some Type 2 objects are shown to be obscured or very faint accretion-powered sources, while others may prove to be completely unrelated to the AGN phenomenon.  This has important ramifications for AGN demographics, because the population of LINER 2s and transition objects outnumbers all other types of AGNs combined.  The fraction of galaxies containing genuine AGNs also sets a lower limit to the fraction of galaxies containing supermassive black holes.  This is still a pertinent issue because dynamical measurements of black hole masses have been performed only for a few dozen galaxies, while AGN surveys can indirectly detect black holes in vastly more galaxies, and over a broader range of Hubble types.  In applying unification ideas to LLAGNs, several questions arise.  To what extent is our view of LLAGNs determined by orientation and obscuration?  Do Type 2 LLAGNs show evidence for obscuring tori in the form of polarized broad-line emission, heavily absorbed X-ray sources, or ionization cones?  And are some objects classified as LLAGNs actually powered predominantly or entirely by stellar processes, rather than by accretion?  Other aspects of the LLAGN phenomenon will be covered in more detail in other contributions to this volume, including those by Nagar (radio observations), Mushotzky (X-ray observations), and Ho (spectral energy distributions and central engine physics).  Additional recent reviews of the properties of LLAGNs, including discussion of unification issues, are given by V\\'eron-Cetty \\& V\\'eron (2000) and by Ho (2002). ",
        "conclusions": "Traditional unification models based on the orientation of an optically thick obscuring torus apparently do apply to at least some fraction of the LLAGN population, as shown by a few detections of emission-line polarization, ionization cones, and heavily obscured X-ray sources.  These detections constitute a small minority of the Type 2 LLAGN population, however.  Obscuration on the scale of the host galaxy plays a more important role in affecting our view of the central regions of LINERs, at least in the UV and optical.  The larger question is whether Type 2 LLAGNs contain genuine AGN-like central engines at all.  Radio and X-ray surveys suggest that the answer is ``yes'' for many, and perhaps most, LINER 2s.  From the results of their \\emph{Chandra} survey, Ho et al.\\ (2001) estimate that at least 60\\% of all LINERs contain AGNs, but this result is still based on very small-number statistics.  Further observations with \\emph{Chandra} will be the best way to improve on this estimate and complete the census of nearby AGNs, both by surveying larger samples and by obtaining deeper exposures to search for obscured or intrinsically very faint sources.  Transition objects as a class present the weakest case for being members of the AGN family, and their spectra may instead be a manifestation of star formation in the extreme environment of galaxy centers."
    },
    "0201/astro-ph0201253_arXiv.txt": {
        "abstract": "We analyze the structural parameters of the largest-available sample of spatially resolved extragalactic globular clusters. The images of M31 GCs were found in a search of HST archival data, described in a companion paper. We measure the ellipticities and position angles of the clusters and conclude that the ellipticities are consistent with being caused by rotation. We find that most clusters' surface brightness distributions are well-fit by two-dimensional single-mass Michie-King models. A few clusters show possible power-law distributions characteristic of core-collapse, but the spatial resolution is not high enough to make definitive claims. As has been found for other galaxies, the metal-rich clusters are slightly smaller than the metal-poor clusters. There are strong correlations between structural properties of M31 GCs, as for Milky Way clusters, and the two populations are located close to the same `fundamental plane' in parameter space. ",
        "introduction": "Globular clusters' structures yield important information about  their dynamical states and also about the environmental effects of the parent galaxy's tidal field. The core radii and ellipticities of Galactic globulars are known to vary with position \\citep{vdb94,ws87}, and the mean structural parameters of clusters also vary between galaxies. For example, the LMC's globular clusters are, on average, much more flattened than those of the Milky Way \\citep{vdb83}. These findings have led to many suggestions and theories about the histories of GCSs and their parent galaxies \\citep[e.g.,][]{vdb00}, but these conclusions would be much stronger if structural parameters were accurately measured for clusters in more galaxies.  Globular clusters in Local Group galaxies are particularly valuable for such a comparison. They are distant enough that their integrated properties can be easily derived, but near enough that their individual stars can also be resolved. M31 has the Local Group's largest globular cluster population, so it is a natural starting place for studies of extragalactic globular clusters. The M31 GCS is known to be similar to the Milky Way's in numerous ways \\citep[e.g., metallicity and spatial distributions;][]{b00}, but has some important differences as well \\citep*[chemical composition and possibly age and luminosity distribution:][]{bh91,bhb01}. While many properties of the M31 globular clusters can be elicited with ground-based imaging and spectroscopy, high spatial resolution data from space-based telescopes such as HST are required to study their spatial structures and stellar populations. The angular sizes of M31 globular clusters are comparable to the sizes of ground-based seeing disks: \\citet{s79,s81} and \\citet{hol_phd} showed that seeing can lead to substantial errors in the derived structural parameters of galaxies with similar sizes. The first HST measurements \\citep{ben93,fp94} showed that the ground-based measurements of core and half-light radii were systematically overestimated. To date, post-repair HST measurements of M31 GC parameters (\\citealp{cmd4,gri96}; \\citealp*{hfr97}) have included only a total of seven objects.  In a companion paper \\citep[hereafter Paper I]{hst_1}, we described the use of HST/WFPC2 images to identify globular clusters and candidates in M31 and thus examine the quality of existing cluster catalogs. That paper contains the details of the HST Archive search and data reduction procedures. In this paper, we derive structural parameters for `good' (category A and B) cluster candidates found in our survey. The objects studied here do not constitute a complete or magnitude-limited sample of globular clusters in M31, and only the clusters which were not specific HST targets can be considered to be a random sample. Some of our new unconfirmed cluster candidates may not be M31 globulars at all Several of the clusters in our sample --- the targets of HST program GO-6699 --- may actually belong to M31's companion galaxy NGC~205. We retain these objects in our sample except when computing correlations between cluster properties and location in the galaxy. We have studied a much larger sample of M31 globular clusters than has been previously available, and this allows us to explore correlations of structural parameters with other properties such as galactocentric distance, and to determine mean structural properties for comparison with other galaxies. ",
        "conclusions": "We use HST images of M31 globular clusters to measure the clusters' sizes, shapes, and best-fit King model parameters. Cluster departures from sphericity are consistent with being caused by rotation, although there are also indications of relations between ellipticity and luminosity and metallicity. We find a slight difference between the half-light radii of metal-rich and metal-poor clusters, consistent with previous results on clusters in other galaxies. The M31 clusters are well-described by the three-parameter family of King models. They have approximately the same range of parameter values as the Milky Way clusters, except that there are few faint, low-concentration clusters in our M31 sample due to observational selection effects. The scatter about the fundamental plane relations is very similar for Milky Way and M31 clusters, although the M31 clusters are offset from the Milky Way cluster relations. This effect may be due to an intrinsic difference in the two galaxies' clusters' fundamental planes or to systematic errors in our measurements: further information is needed. The overall similarity of the two fundamental planes implies that the formation and evolution of GCs must have been very similar in the two galaxies."
    },
    "0201/astro-ph0201123_arXiv.txt": {
        "abstract": "We investigate the predicted Color-Magnitude distribution of metal-rich Horizontal Branch (HB) stars, discussing selected theoretical models computed under various assumptions about the star metallicity and the efficiency of super-adiabatic convection.  We find that canonical Zero Age Horizontal Branches with metallicity larger or of the order of Z=0.002 should be all affected by a tilt, by an amount which increases when the metallicity is increased and/or the mixing length is decreased, reaching a tilt of $\\Delta V \\sim$0.2 mag in the case of solar metallicity when a mixing length value $\\alpha$=1.6 is assumed ($\\Delta V$ is the magnitude difference between the top of the blue HB and the fainter magnitude reached by the red HB). Uncertainties in the luminosity of the red HB due to uncertainty in the mixing length value are discussed.  We finally discuss the much larger tilt observed in the clusters NGC~6441 and NGC 6388, reporting additional evidence against suggested non-canonical evolutionary scenarios.  Numerical experiments show that differential reddening could produce such sloped HBs.  Further, HST-PC imaging of NGC~6441 gives clear indications about the occurrence of differential reddening across the cluster. However, the same imaging shows that the observed slope of the red HB {\\em is not} an artifact of differential reddening.  We finally show that sloping red HBs in metal rich clusters are a common occurrence not necessarily correlated with the appearance of extended blue HB. ",
        "introduction": "In a recent paper Brocato et al. (1999, paper I) we have presented the case of the intermediate metallicity globular cluster NGC 6362, where Horizontal Branch (HB) stars reach a maximum V luminosity at (B$-$V)$\\sim$0.2 mag, with an overall tilt of the order of $\\Delta V$$\\sim$0.1 mag, thus of the same order of magnitude of the tilted HB already observed in NGC 1851 (Walker 1998), possibly in NGC 6229 (Borissova et al. 1999) and in NGC 6712 (Ortolani et al. 2000).  In the same paper it has been shown that canonical predictions concerning low-mass, He-burning stars with Z= 0.002 do foresee such an occurrence, nicely accounting for the observed NGC 6362 HB stars distribution. Thus we regard such a tilt as an evolutionary feature which should characterize all the well populated intermediate metallicity clusters when observed with sufficient photometric accuracy.  The evidence for such a \"canonical\" tilt, as due to the behavior of the bolometric correction, appears of some relevance in connection with the debated problem of the much larger tilts ($\\Delta V$$\\sim$0.5 mag) observed in some metal rich globulars of the inner Milky Way. The first evidence for similar unexpectedly large tilts was brought to the light from HST observations of the globulars NGC~6388 and NGC~6441 (Piotto et al. 1997, Rich et al. 1997). Sweigart and Catelan (1998) drove the attention on such a feature, regarding it as an observational evidence requiring non-canonical evolutionary scenarios. In the present paper we investigate evolutionary predictions concerning the Color-Magnitude (CM) distribution of metal rich HB stars. In the next section we will present new sets of canonical HB models, whereas section 3 will deal with their CM diagram and theoretical uncertainties, due to the efficiency of the convection in superadiabatic layers, which affect the predicted luminosity of stars at the red HB end.  In section 4 we will provide arguments and constraints to understand the origin of the large tilt of the high metallicity clusters NGC~6441 and NGC~6388.  Final remarks will close the paper.   \\section {Metal-rich Zero Age Horizontal Branch models}  The recent literature has already devoted considerable attention to metal rich HB structures. Starting from the pioneering paper by Horch, Demarque \\& Pinsonneault (1992) several works have been presented covering this issue with rather detailed evaluations (Dorman, Rood \\& O'Connel 1993, Fagotto et al. 1994, Yi et al. 1997, Bono et al. 1997, Girardi et al. 2000). However, these studies have been mainly devoted to investigating the evolution of HB structures in connection with either the UV excess in elliptical galaxies or the pulsational properties of metal rich RR Lyrae variables and, to our knowledge, up to recent times no attention was devoted to the predicted distribution of metal rich HB stars in the CM diagram.  To investigate such an issue, we decided to extend the evolutionary computations given in Paper I for Z=0.002 to higher metallicities.  Similar computations have been recently presented by VandenBerg et al. (2000) but for a more restricted range of metallicity, not addressing the problem of tilt.  However, this will allow an useful comparison between the results of the two theoretical scenarios for a common value of Z.   In this paper, Zero Age Horizontal Branch (ZAHB) models have been produced by following selected stellar models all along their H-burning phase till the ignition of the He flash, deriving in this way the mass of the He core of new born He burning stars with an age of the order of 10 Gyr, in agreement with several recent estimates of the age of metal rich GCs (Salaris \\& Weiss 1997, Gratton et al. 1997, Chaboyer et al. 1998, Cassisi et al. 1999).  For a detailed discussion on the physical inputs adopted in the computations we refer to Cassisi et al. (1998).  Table~1 gives selected parameters for the above quoted models.  Left to right one finds: original He abundance, metallicity and mass of the model, mass of the He core at the He ignition, surface He abundance after the first dredge up and, finally, star age and luminosity at the RGB tip.  As shown in Table 1, to extend the investigation in Paper I at higher metallicities, we selected two \"suitable\" values of Z, choosing Z=0.006 as representative of \"metal rich\" clusters like 47~Tuc and Z=0.02 as a safe upper limit for globular cluster metallicities.  As in Paper I, the present computations have been performed by adopting a mixing length parameter $\\alpha$=1.6, implemented with the case $\\alpha$=1.0 to explore the effect of mixing length assumptions. Figure~1 gives the predicted distribution of these models in the log(L/L$_{\\odot})$, logT$_{e}$ HR diagram. As labelled in the figure, the amount of original He (Y$_{MS}$) has been increased with Z according to $\\delta$Y/$\\delta$Z $\\sim$2.5, as suggested by current evaluation on that matter (see, e.g., Peimbert 1995, Fernandes et al. 1998).  As a whole, one finds that data in both Table~1 and Figure~1 appear in general agreement with previous evaluations given by Dorman, Rood \\& O'Connel (1993) and Yi et al. (1997), our models being slightly brighter as a consequence of larger initial He cores (as discussed in Cassisi et al. 1998).  As expected, Figure~1 shows that the effective temperature of red HB stars depends on the assumption about the mixing length. This implies that, at any given temperature, the luminosity of the red portion of the ZAHB depends on the assumptions on the mixing length, increasing when the mixing length increases.  This dependence grows with the stellar metallicity, since at higher metallicities the HB locus moves toward lower effective temperatures and red HB star are increasingly affected by external convection, making stellar models more sensitive to the adopted efficiency of the super-adiabatic convection. As we will discuss in the next section, this will leave some uncertainty in theoretical predictions on tilted HB. ",
        "conclusions": ""
    },
    "0201/astro-ph0201315_arXiv.txt": {
        "abstract": "{ In this article I review the dominant physical mechanisms in the extended emission line regions (EELR) of powerful radio galaxies at low redshifts and, in particular, the balance between jet-induced shocks and AGN photoionization. I consider the evidence for jet-induced shocks based on morphological, kinematical and ionization (diagnostic diagram) information. Although each of these three types of information separately provides some evidence for the effects of jet-induced shocks, this evidence is often ambiguous. The major advance in recent years has been to combine morphological, kinematical and diagnostic diagram approaches. In this way it has been possible to show that jet-induced shocks dominate the ionization and acceleration of the EELR along the radio axes of many powerful radio galaxies. However, AGN photoionization is likely to remain the dominant mechanism in the nuclear regions of radio galaxies in which the radio sources extend well beyond the emission line regions.} \\listofauthors{C.~N.~Tadhunter} \\indexauthor{Tadhunter, C.~N.}  \\begin{document} ",
        "introduction": "\\label{sec:intro} The extended emission line regions (EELR: $1 < r < 100$kpc) around powerful radio galaxies have the potential to provide key information about the nature of the central energy sources (e.g. \\pcite{rawlings91}), the triggering and fuelling of the activity (e.g. \\pcite{tadhunter89c}, \\pcite{baum92}), and the evolution of the host galaxies. However, if we are to use them in this way, it is crucial to understand the dominant physical mechanisms. In particular: to what extent are the observed emission line features a consequence of the activity, and to what extent do they reflect the intrisic (pre-activity) properties of the host early-type galaxies?  In this review I will concentrate on the issue of the ionization mechanism for the emission line gas, which has been a perrenial problem for all types of active galaxies since they first started to be studied in depth more than 50 years ago. Like so many fields of astrophysics, this is an area which has benefitted from the considerable improvements in instrumentation and detector technology in the last two decades. At optical wavelengths is now possible to combine high spectral resolution and sensitivity, with wide spectral and spatial coverage. Reflecting this, most of the recent advances in this field have come about by combining ionization, kinematic and spatial information.  I will start by reviewing the main models that have been proposed to explain the extended emission line properties of powerful radio galaxies at low redshifts ($z < 0.4$). I will then consider separately the approaches of using morphological, kinematical and emission line ratio information to determine the dominant ionization mechanism, before demonstrating the power of combining these approaches. ",
        "conclusions": "\\label{sec:discuss}  The emission line nebulae around powerful radio galaxies are complex, but by using a combination of morphological, kinematical and ionization information it is now possible to start to distinguish the dominant physical mechanisms. The strongest result is that there is now compelling evidence that at least a subset of EELR aligned along the radio axes of powerful radio galaxies are ionized and accelerated in jet-induced shocks. The main outstanding questions are as follows. \\begin{itemize} \\item What is the balance between AGN photoionization and shock ionization (including shock-photoionized precursor) for the EELR that are not obviously associated with radio features? \\item Why do powerful radio galaxies at high redshifts ($z > 0.6$) appear to show stronger evidence for jet-induced shocks than their lower redshift counterparts? \\item Given that the broad wings to the emission lines are likely to be symptomatic of the entrainment and eventual destruction of the clouds in the host post-shock wind \\cite{villarmartin99}, is jet-induced star formation viable? \\item How important are starburst- or AGN-driven winds in the near-nuclear regions of powerful radio galaxies? \\end{itemize} Key future observations are likely to include: integral field spectroscopy to accurately measure the velocity shear and ionization of the shocked gas relative to the ambient ISM; and deep X-ray imaging and optical spectroscopy to study the hot cooling gas and thereby determine  accurate shock parameters."
    },
    "0201/astro-ph0201301_arXiv.txt": {
        "abstract": "On the average, 1.5 new publications on cosmic gamma-ray bursts enter the literature every day.  The total number now exceeds 5300. I describe here a relatively complete bibliography which is on the web, and which can be made available electronically in various formats. ",
        "introduction": "I have been tracking the gamma-ray burst literature for about the past twenty-one years, keeping the authors, titles, references, and key subject words in a machine-readable file.  The present version updates previous ones reported in 1993 \\cite{hurley1}, 1995 \\cite{hurley2},1997 \\cite{hurley3}and 1999 \\cite{hurley4}. In its current form, this information is in a Microsoft Word 97 \"doc\" format.  My purpose in doing this was first, to be able to retrieve rapidly any articles on a given topic, and second, to be able to cut and paste references into manuscripts in preparation.  The following journals have been scanned on a more or less regular basis starting with the 1973 issues:\\\\ \\linebreak Advances in Physics\\\\ Annals of Physics\\\\ Astronomical Journal\\\\ Astronomische Nachrichten\\\\ Astronomy and Astrophysics (including Supplement Series)\\\\ Astronomy and Astrophysics Review\\\\ Astronomy Letters (formerly Soviet Astronomy Letters)\\\\ Astronomy Reports (formerly Soviet Astronomy)\\\\ Astrophysical Journal (letters, main journal, and supplements)\\\\ Astrophysical Letters and Communications\\\\ Astrophysics and Space Science\\\\ ESA Bulletin\\\\ ESA Journal\\\\ IAU Circulars\\\\ IEEE Transactions on Nuclear Science\\\\ Journal of Astrophysics and Astronomy\\\\ Monthly Notices of the Royal Astronomical Society\\\\ Nature\\\\ Nuclear Instruments and Methods in Physics Research Section A\\\\ Observatory\\\\ Physical Review (main journal A and letters)\\\\ Proceedings of the Astronomical Society of Australia\\\\ Publications of the Astronomical Society of Japan\\\\ Publications of the Astronomical Society of the Pacific\\\\ Reports on Progress in Physics\\\\ Science\\\\ Scientific American\\\\ Sky \\& Telescope\\\\  In addition, the following journals either have been scanned, but less regularly in the past, or in some cases, are no longer being scanned:\\\\ \\linebreak Annals of Geophysics\\\\ Astrofizika\\\\ Astroparticle Physics\\\\ Bulletin of the American Astronomical Society\\\\ Bulletin of the American Physical Society\\\\ Bulletin of the Astronomical Society of India\\\\ Chinese Astronomy and Astrophysics\\\\ Chinese Physics Letters\\\\ Cosmic Research\\\\ Journal of Atmospheric and Terrestrial Physics\\\\ Journal of the British Interplanetary Society\\\\ Journal of the Royal Astronomical Society of Canada\\\\ New Astronomy\\\\ Progress in Theoretical Physics\\\\ Solar Physics\\\\ Soviet Physics\\\\  The above lists are not exhaustive.  For example, where theses or internal reports have come to my attention, I have included them, too.  To be included, an article had to have something to do with GRB or SGR  theory, observation, or instrumentation, or be closely related to one of these topics (e.g., merging neutron stars, AXPs, high-z supernovae, etc.), and must have been published.  With only a few exceptions, preprints or internal reports which were never published have not been included. ",
        "conclusions": ""
    },
    "0201/astro-ph0201137_arXiv.txt": {
        "abstract": "We describe the BOOMERanG experiment and its main result, i.e. the measurement of the large scale curvature of the Universe. BOOMERanG is a balloon-borne microwave telescope with sensitive cryogenic detectors. BOOMERanG has measured the angular distribution of the Cosmic Microwave Background on $\\sim 3\\%$ of the sky, with a resolution of $\\sim 10$ arcmin and a sensitivity of $\\sim 20 \\mu K$ per pixel. The resulting image is dominated by hot and cold spots with rms fluctuations $\\sim 80 \\mu K$ and typical size of $\\sim 1^o$. The detailed angular power spectrum of the image features three peaks and two dips at $\\ell = (213^{+10}_{-13}), (541^{+20}_{-32}), (845^{+12}_{-25}  )$ and $\\ell = (416^{+22}_{-12}), (750^{+20}_{-750})$, respectively. Such very characteristic spectrum can be explained assuming that the detected structures are the result of acoustic oscillations in the primeval plasma. In this framework, the measured pattern constrains the density parameter $\\Omega$ to be $0.85 < \\Omega < 1.1$ (95\\% confidence interval). Other cosmological parameters, like the spectral index of initial density fluctuations, the density parameter for baryons, dark matter and dark energy, are detected or constrained by the BOOMERanG measurements and by other recent CMB anisotropy experiments. When combined with other cosmological observations, these results depict a new, consistent, cosmological scenario. ",
        "introduction": "The almost isotropic Cosmic Microwave Background (CMB) accounts for most of the photons present in our Universe. These photons have been produced in the very early  Universe, and were last scattered by free electrons at recombination, about 300000 years after the Big Bang. After that, CMB photons travel basically undisturbed along spacetime geodesics for $\\sim$ 15 Gyr, reaching our telescopes redshifted by a factor $\\sim 1000$ due to the expansion of the universe.  Thus, when we look to the image of the CMB, we see the result of early processes (e.g. the generation of density perturbations at $t \\sim 10^{-30} s$ after the Big Bang, and matter-antimatter annihilations, at $t \\sim 10^{-5}$ s), of processes in the plasma era till recombination (e.g. acoustic oscillations of the matter-photons plasma till $t \\sim 300000 yrs$), and of the large scale geometry of the universe (affecting the geodesics followed by the photons after recombination, from $t \\sim 300000 yrs$ to $t \\sim 15 Gyrs$ ). Any small curvature of the Universe and of the geodesics would affect significantly the image of the CMB, magnifying or demagnifying it with respect to the Euclidean case. For this reason experiments mapping the CMB are very sensitive to the curvature of the Universe, which, according to General Relativity, is determined by the average mass-energy density of the Universe $\\rho$.  A parametric, general relativistic theory of the anisotropy of the CMB has been fully developed in the last 35 years, based on the main cosmological observations: the isotropic expansion of the Universe, the primordial abundances of light elements, the existence of the CMB and its black-body spectrum, the large scale distribution of Galaxies. Detailed models and codes are available to compute the angular power spectrum of the CMB image given a cosmological model for the generation of density fluctuations in the Universe, and a set of parameters describing the background cosmology \\cite{Hu}.  The default model is nowadays the inflationary adiabatic one, where gaussian, adiabatic density fluctuations are generated from quantum fluctuations of a scalar field present in the very early Universe, boosted to cosmological scales by the inflation phase at $E \\sim 10^{15} GeV$ \\cite{KT1990}. This modification of the standard theory is needed in order to solve several paradoxes intrinsic to the standard Hot Big Bang theory. As we will see below, this theory offers a natural explanation to several cosmological observations, but the ingredients required to fit the data are non trivial, requiring the presence of unobserved \"dark matter\" and unknown \"dark energy\". \\begin{figure}[tb] \\begin{center} \\begin{minipage}[t]{16 cm} \\epsfig{file=fig1.eps,scale=1.0} \\end{minipage} \\begin{minipage}[t]{16.5 cm} \\caption{Position of the first peak in the angular power spectrum of the CMB as a function of the total matter-energy density parameter $\\Omega$. The results are plotted for a wide database of adiabatic inflationary models ($\\sim 5\\times 10^5$ models corresponding to the ranges of the parameters listed in the bottom panel), for three sample values of $\\Omega_\\Lambda$ (as in the top panel). The thick lines are the average values of $\\ell_1$. The band within thin lines includes all the possible combinations of the other parameters. See text for details and use of the graph. \\label{fig1}} \\end{minipage} \\end{center} \\end{figure}  The parameters of the model are: the Hubble constant $H_o$ (setting the expansion rate of the Universe); the density parameter $\\Omega$ (i.e. the ratio between $\\rho$ and the critical density $\\rho_c = 3H_o^2 / 8 \\pi G $); the baryons, dark matter and dark energy density parameters $\\Omega_b$, $\\Omega_m$, $\\Omega_\\Lambda$; the spectrum of the primordial density fluctuations generating CMB anisotropy. The latter is usually expressed as a power law with spectral index $n$.  The general prediction of the model above is a series of \"acoustic peaks\" in the angular power spectrum of the CMB. These peaks derive from acoustic oscillations of the photons-matter plasma (the primeval fireball).  Density perturbations $\\Delta \\rho / \\rho$ were oscillating in the primeval fireball as a result of the opposite effects of the pressure of photons and of gravity. After recombination, photons pressure becomes unimportant, and $\\Delta \\rho / \\rho$ can grow and create, through gravitational instability, the hierarchy of structures we see today in the nearby Universe. There are three physical processes converting the density perturbations present at recombination into {\\it observable } CMB temperature fluctuations $\\Delta T / T$. They are: the photon density fluctuations $\\delta_\\gamma$, which can be related to the matter density fluctuations $\\Delta \\rho$ once a specific class of perturbations is specified; the gravitational redshift of photons scattered in an over-density or an under-density  with gravitational potential difference $\\phi_r$; the Doppler effect produced by the proper motion with velocity $v$ of the electrons scattering the CMB photons. In formulas: \\begin{equation} \\frac{\\Delta T}{T}(\\vec{n}) \\approx \\frac{1}{4}\\delta_{\\gamma r}+ \\frac{1}{3} {\\phi_r \\over c^2} - \\vec{n} {\\vec{v_r} \\over c} \\end{equation} where $\\vec{n}$ is the line of sight vector and the subscript $r$ labels quantities at recombination.  There is an acoustic horizon at recombination, with a linear size of $\\sim 300000$ light-years. At any time $t$ in an infinite Universe regions separated more than $c t$ are not in causal contact. This means that causal horizons, with size $ct$, exist in the Universe. The acoustic horizon has a size similar to the causal horizon, since in the primeval plasma the speed of sound is close to the speed of light. Since the horizon expands with time, any given proper size will eventually become smaller than the acoustic horizon. The time when the proper size of a perturbation becomes equal to the acoustic horizon is called horizon crossing. Perturbations larger than the acoustic horizon at recombination have never been in acoustic contact before, and have not been able to oscillate. Perturbations smaller that that size have oscillated after crossing the horizon, and arrive to recombination with a phase depending on their intrinsic size. The image of the CMB, which is directly observable by means of CMB anisotropy experiments, is a processed image of density perturbations present at recombination. The size of the acoustic horizon is expected to be evident in the size distribution of the detected CMB structures $\\Delta T(\\alpha, \\delta)$. In fact, perturbations with a size close to the acoustic horizon had just enough time to fully compress or rarefy before recombination, and will be evident as cold or hot spots. If we compute the angular power spectrum $c_\\ell$ of the image, where $c_\\ell = \\langle |a_{\\ell, m}|^2 \\rangle$, and $\\Delta T(\\alpha, \\delta) = \\sum_{\\ell, m} a_{\\ell, m} Y^\\ell_m (\\alpha, \\delta)$, we expect to see a peak corresponding to the angular distance subtended by the acoustic horizon at recombination. In an Euclidean universe, the subtended angle is simply 300000 lyr divided by 15 Glyr and multiplied by a factor 1000, to take into account the subsequent expansion of the Universe. We thus expect structures with a typical angular scale $\\theta_1 \\sim 1^o$, which corresponds to a multipole $\\ell_1 \\sim \\pi / \\theta_1 \\sim 200$.  This is correct only if the geometry of the Universe is Euclidean, not curved, i.e. if the average mass-energy density of the Universe is the critical one ($\\Omega=1$). If, instead, the mass-energy density is higher than critical ($\\Omega > 1$), the geometry of space will have a positive curvature and the photons will travel along curved geodesics. The excess density will act as a magnifying glass, and the same fluctuations in the CMB will appear as spots larger than $1^o$.  The opposite will happen if the density is lower than critical, acting as a de-magnifying glass and producing a typical angular size of the fluctuations smaller than $1^o$. By measuring the location of the peak it will thus be possible to measure $\\Omega$. The quantitative treatment of this angular-size vs distance test can be found in \\cite{Wein2000}, \\cite{Mel00}. In general, $\\ell_1$ decreases when $\\Omega$ increases, but the location of the first peak is also controlled by $\\Omega_\\Lambda$, which effectively changes our distance from recombination (see fig.1). Only if $\\Omega_\\Lambda =0$ the simple relationship $\\ell_1 \\sim \\Omega^{- {1\\over 2}}$ holds \\cite{Wein2000}.  CMB experiments are now starting to measure the angular power spectrum with  sufficient accuracy to infer $\\Omega$ and several other cosmological parameters. This inverse procedure, i.e. measure the cosmological parameters given the observed angular power spectrum and its measurement error, is a complex one, and is further complicated by the presence of degeneracies between the cosmological parameters. Different combinations of the parameters can generate the same power spectrum of the CMB \\cite{Efs99}, so priors (coming from independent cosmological evidence) must be assumed in the analysis \\cite{Lan01}. The measurement of curvature is quite robust in this sense. In fig.1 it is evident that a measurement of e.g. $\\ell_1 = 200$ is consistent only with models with $0.85<\\Omega<1.1$, for any possible value of the other parameters $\\Omega_\\Lambda, \\Omega_b, \\Omega_m, h, n$. More precise measurements of $\\Omega$ can be obtained from a full likelihood analysis of the power spectrum data (see below). Once the default model and rather weak priors (for example on the Hubble constant) are assumed, the measurements set very strong constraints also on all the other parameters, and the results are consistent with independent cosmological observations (see e.g. \\cite{Teg01}. A new \"precision phase\" in the cosmological research seems to be starting. This, however, has a cost, which is the introduction of \"strange\" components in the Universe like dark matter and dark energy. Only further observations will show if the latter are just artifices, like Ptolemy's epicycles deferents and late additions, or, instead, are very important new discoveries.  In this paper we show how the curvature of the Universe has been measured by BOOMERanG. We describe the experiment, then the experimental strategy and observations obtained in the long duration flight, and finally the cosmological implications. ",
        "conclusions": "The BOOMERanG experiment has produced multi-frequency maps of the microwave sky, where the structure of the CMB has been resolved with high signal to noise ratio. The structures in the CMB are gaussian, and their power spectrum features three peaks. This is consistent with the presence of acoustic oscillations in the primeval plasma. It also fits the predictions of the adiabatic inflationary scenario. The values of the cosmological parameters inferred in this scenario point to a flat universe (the 95\\% confidence interval is $0.85 < \\Omega < 1.1$) with nearly scale-invariant initial adiabatic perturbations and a significant contribution of dark energy to the total density of the Universe. These results from BOOMERanG have been confirmed by independent CMB experiments (like DASI and MAXIMA) and by other cosmological observations. The forthcoming space missions MAP and Planck will improve significantly the precision of these results, either entering the \"precision cosmology\" era, or detecting hidden inconsistencies of the present cosmological scenario."
    },
    "0201/astro-ph0201247_arXiv.txt": {
        "abstract": "{ Following on from the numerical work of Capelato, de Carvalho \\& Carlberg (1995, 1997), where dissipationless merger simulations were shown to reproduce the ``Fundamental Plane\" (FP) of elliptical galaxies, we investigate whether the end products of pure, spherically symmetric, one-component dissipationless {\\it collapses} could also reproduce the FP.  Past numerical work on collisionless collapses have addressed important issues on the dynamical/structural characteristics of collapsed equilibrium systems. However, the study of collisionless collapse in the context of the nature of the FP has not been satisfactorily addressed yet. Our aim in this paper is to focus our attention on the resulting collapse of simple one-component spherical models with a range of different initial virial coefficients.  We find that the characteristic correlations of the models are compatible with virialized, centrally homologous systems. Our results strengthen the idea that merging may be a fundamental ingredient in forming non-homologous objects. ",
        "introduction": "Self-gravitating stellar systems, ranging from globular clusters to clusters of galaxies, show significant correlations among kinematic and photometric parameters (\\cite {bur97}).  These correlations are important tracers of the formation histories of these structures. However, the origin of these scale relations are not yet known clearly.  For instance, the ``Fundamental Plane'' (FP, c.f. \\cite{djo87}, \\cite{dre87}) relation of elliptical galaxies represent a significant departure from the prediction of the virial theorem, under the assumption that ellipticals are simple, one-component, homologous systems.  This relation is described by: $r_e \\sim \\sigma_0^A I^B$ (where $\\sigma_0$ is the central velocity dispersion, $I$, the average surface brightness within the effective radius in linear units, and $r_e$ is the effective radius, where $A \\sim 1.53$, $B \\sim -0.79$ (e.g. \\cite{pah98}).  One current hypothesis to explain the observed discrepancy postulates that the mass-luminosity ratio of ellipticals would be a function of total luminosity (e.g. \\cite{djo93}, \\cite{djo88}, \\cite{pah95}). An alternative hypothesis takes into consideration the possible effects of the dark matter halo on the FP correlations (\\cite{dan00}). Another explanation (e.g., \\cite{hjo95}) is based on the assumption that the homology hypothesis is not valid, so that elliptical galaxies would be non-homologous virialized systems.  Several works have addressed the latter possibility (e.g., \\cite{cio96}, \\cite{bus97}, \\cite{gra97}, \\cite{bek98}).  In particular Capelato, de Carvalho \\& Carlberg (1995, 1997, hereon CdCC95 and CdCC97) showed that the FP correlations arise naturally from objects that are formed by dissipationless hierarchical mergers of pre-existing galaxies.  The end product of their simulations was a non-homologous family of objects following almost exactly the observed K-band FP, with a scatter that was only half of that observed.  Merging is a natural process in a hierarchical galaxy formation scenario, and the observed structural properties of ellipticals seem to be well accounted for by this mechanism (e.g. \\cite{shi98}, \\cite{ben99} and references therein).  On the other hand, theoretical and numerical investigations on pure dissipationless collapses of stellar systems have historically been of special interest (\\cite{pol81}, \\cite{hen73}, \\cite{alb82}, \\cite{mcg84}, \\cite{vil84}, \\cite{may84}, \\cite{mer85}, \\cite{agu90}, \\cite{lon91}, \\cite{kat91}, \\cite{car95}). In particular, the role of gravitational instabilities in a model evolving towards equilibrium is still an open field of investigation (\\cite{pal94}), and and may be a key feature for understanding the basics of the dynamics that play a role in the formation of galaxies.  However, previous numerical works have not fully addressed the dynamical/structural characteristics of collapsed systems in the context of the FP. It is also a question for investigation whether the results of the merger simulations by CdCC95 can be reproduced by other initial conditions, that is, whether the FP could arise solely as a function of the dynamics of relaxation. A preliminary discussion of these subjects has been given in CdCC97 where it is suggested that, under appropriate initial conditions, dissipationless collapses could also follow a FP relationship. We believe that a first move towards answering the questions raised above should focus on the simplest dynamical conditions. For this reason, we restrict our analysis to the classical scenario of elliptical galaxy formation through one-component dissipationless gravitational collapse.  Evidently, the effects of a second component (extended dark halo) on the final equilibrium conditions of the luminous matter, as much as the presence of gas dissipation, should be considered in a more refined analysis.  This paper is organized as follows: in Section 2, we present the simulations and define the characteristic parameters; in Section 3, the end products of the simulations are considered in the context of the FP analysis, along with their dynamical/structural properties. In section 4, we discuss our results. ",
        "conclusions": "We designed a set of numerical experiments of dissipationless collapse of spherical models in order to investigate some possible constraints on the origin of the FP of elliptical galaxies.  First, the final objects are found at different locations on the FP parameter space, depending on the initial model. Even if the initial models have the same mass and length (e.g. $R$) scales, which is the case for the K and A models, they will occupy almost non-intersecting regions on the FP space. The central mass concentration of the initial models may have a determining influence on the final locations of the models on the FP space, for the same set of global structural-dynamical parameters (viz. $\\beta$ and $R$).  Second, we found that the collapsed models do not globally populate the FP, being more compatible with homologous virialized systems.  In particular, the final ``coldest'' C models, especially the n=1 models, tend to form a family with a deviation from homology, although the error bars are too large to confirm this effect. The C models are not only the result of a global collapse of the system, but also the product of a series of mergers of small fluctuations of smaller scales, which had the opportunity to grow depending on the initial conditions. In particular, the $n=0$ models seem to have collapsed more homogeneously, since its resulting scaling relations slope is quite similar to the K and A models, which are initially spherically symmetric objects with diferent central mass concentrations. In fact, the $n=0$ models represent a homogeneous model where small-scale noise has been added, so that sufficiently large clumps probably were not formed. This is unlike the $n=1$ and $n=2$ models, that might have resulted from mergers of significant clumps along with the main global collapse.  From the above discussion and also from the results of CdCC95, we conclude that merging may indeed be an important ingredient in forming non-homologous objects. In addition, we consider that, although ``hot'' collapses result in a ``softer'' process toward virialization (so that the relaxation mechanism of these collapses should be more closely related to the one produced by a slow merger), it is the mechanism of merging {\\it per se} that drives objects towards non-homology. This is because we found that ``cold'' collapses from initial clumpy conditions are non-homologous, whereas the same cold collapse factors applied to more spherically symmetric systems {\\it do} produce homology. From this, we also deduce that any dynamical constraints dictated by spherical symmetry in the initial conditions are not likely to have been present in the formation of elliptical galaxies, at least for those in clusters.  This result is of course consistent with those found in the classical literature, where reasonably strong indications are found that ``clumpy'' and ``cold'' initial conditions are fundamental for the formation of objects similar to ellipticals (e.g.  \\cite{alb82}, \\cite{may84}, \\cite{mer85}, \\cite{agu90}, \\cite{lon91}).  \\begin{figure*} \\centering \\includegraphics[width=12cm]{MS1622f4.eps} \\caption{Field ellipticals data from de Carvalho \\& Djorgovski (1992) in terms of the FP parameters. The solid line is a homologous fit and the dotted line is a FP fit for the Coma cluster.} \\label{elipcampo} \\end{figure*}  On the other hand, considering the possibility that {\\it field} elliptical galaxies underwent insignificant merging episodes in their past histories (e.g. \\cite{tot98}), and assuming that their gross properties could be described by a simple, spherically symmetric collapse model, our simulations indicate that they should present a FP relation much closer to the virial expectations, compared to cluster ellipticals. Alternatively, these field elliptical galaxies could also be non-homologous objects if they collapsed from reasonably ``clumpy'' and ``cold'' initial conditions. Recently, some authors have concentrated on the analysis of the FP of field early-type galaxies, but the data is still very preliminary and precludes any firm determination of the resulting FP ``tilt'' for these galaxies (e.g.,\\cite{dok01}, \\cite{tre99}, \\cite{pah98}, \\cite{dcd92}).  For instance, we directly tested the above hypothesis using the field ellipticals data from de Carvalho \\& Djorgovski (1992). In Figure \\ref{elipcampo}, we show this data in terms of the FP parameters. The solid line is a homologous fit  and the dotted line is a FP fit for the Coma cluster.  The residuals show that it is not possible to verify which of the relations is the best fit due to the high dispersion of the data.  The K models in our simulations are able to produce some dispersion, depending on the initial parameters (e.g. $\\beta$ and/or mass), but they are not able to ``move'' the resulting models along any relation.  We conclude that much more data is needed in order to decisively test the constraints presented by our models on the FP of field ellipticals.  Finally, we point out that our present analysis is, of course, not entirely realistic, and is based on simple one-component dissipationless models.  Our conclusions should be interpreted more as a {\\it general trend} to be further investigated by a larger series of higher resolution simulations. But the exercise seems to be useful in bringing  out some clues that otherwise could be difficult to unravel in a more complex scenario.  We intend to further refine our results adopting more realistic scenarios. Two-component models and full cosmological simulations are being currently investigated and will be the subject of a future paper."
    },
    "0201/astro-ph0201071_arXiv.txt": {
        "abstract": "We report on Chandra imaging observations of the Galactic Unidentified $\\gamma$-ray source GEV J1809-2327, comparing the X-ray images with new VLA 1.46 GHz and 4.86 GHz maps. The X-ray images reveal a point source connected to a non-thermal X-ray/radio nebula, supporting a pulsar/wind model for the $\\gamma$-ray emitter. We also detect numerous X-ray sources from the young stellar association in the adjacent HII region S32. ",
        "introduction": "The \\EGRET\\ instrument on the {\\it Compton Gamma Ray Observatory} detected many bright sources along the Galactic plane, most of which remain unidentified. A number of these sources have high significance in $E> $GeV photons \\citep{lam97}, providing relatively good source localizations. The key to making progress on identifications is to search for lower energy, particularly hard X-ray and radio, counterparts. \\citet*[RRK]{rob01} have carried out an \\ASCA\\ survey of the bright GeV sources to search for counterparts and test the nature of the GeV emitters. \\bigskip  \\begin{inlinefigure} \\epsscale{0.95} \\plotone{f1.eps} \\figcaption{VLA 21cm map with ACIS 0.5-8keV smoothed contours. The ACIS-I FOV (square) and the brightest sources of Table 1 are also indicated.} \\label{xray-radio} \\end{inlinefigure}  GeV J1809-2327(=2EG J1811-2339) is one of the brightest, best-localized unidentified plane sources. It is one of a subset of sources showing evidence of $\\gamma$-ray variability, with an extended hard-spectrum X-ray counterpart argued to be a pulsar synchrotron wind nebula, or PWN \\citep{rob01}. This source is also interesting because it is adjacent to the dark cloud Lynds 227 and to the HII region S32 with an embedded association of young high-mass stars. These are plausibly associated with the X-ray source and at $\\sim 1.8$kpc provide a possible birth site for a pulsar. \\citet{oka99}, in a study of the molecular gas in this region, found morphological support for this association and hypothesized that the $\\gamma$-ray photons were produced when TeV pulsar electrons penetrated the dense molecular gas, suffering relativistic bremsstrahlung losses.  Studies of this complex region were hampered by the modest angular resolution of the \\ASCA\\ \\GIS\\ ($\\sim 3\\arcmin$ half power diameter), so we have obtained high resolution \\ACIS\\ images to separate out the point source contribution to the X-ray complex. This is abetted by new VLA continuum maps. These data support the pulsar/PWN hypothesis and give new constraints on the origin of the high energy emission. ",
        "conclusions": "Comparing our new observations with the pulsar+PWN hypothesis, we note that source 5 composite spectrum is characteristic of a young pulsar, with a hard (presumably magnetospheric) power law and underlying thermal emission. Given the high $kT=0.18$keV, our fit thermal flux implies a small effective area of $ {\\rm A_{eff}}\\approx 4 \\times 10^{11} (d/1.8{\\rm kpc})^2 {\\rm cm^2}. $ At $\\la 3$\\% of a neutron star's surface, this could represent a polar cap heated by precipitation of magnetospheric particles. Alternatively, atmosphere effects \\citep{ro87,pav01} can produce a hard thermal tail, although such a large fit $kT$ is difficult to produce for reasonable pulsar ages. At a more plausible  $kT_{eff} \\la 0.05\\kev$, full surface emission from cooling would be unobservable; our data only limit such a component to $< 1.2 \\times 10^6$K.  Both the power law emission and the hot thermal component can thus measure magnetospheric activity, and have been phenomenologically related to the spin-down luminosity. \\citet{bt97} find that for pulsars observable by ROSAT, the flux scaled as  $L_x(0.1-2.4\\kev) \\approx 10^{-3} {\\dot E}$; after correcting for absorption our inferred $0.1-2.4\\kev$ flux of $\\approx 2 \\times 10^{-12} \\ecs$ corresponds to a spin-down luminosity ${\\dot E} = 8 \\times 10^{35} d_{1.8}^2\\es$ at the fiducial source distance. \\citet{sai97} find that the (pulsed) \\ASCA\\ luminosity scales as $L(2-10\\kev) \\approx 10^{34} {\\dot E_{38}}^{3/2} \\es$. Using our 2-10keV unabsorbed point source flux, we infer from this relationship ${\\dot E} = 1.5 \\times 10^{36} d_{1.8}^{4/3} \\es$.  For dipole spindown we have ${\\dot E} \\approx 10^{38} (B_{12} \\tau_4)^{-2} \\es$; these luminosities indicate typical spin parameters for a $\\gamma$-ray emitting pulsar of $B_{12} \\tau_4 \\approx 10$, where the surface dipole field is $10^{12}B_{12}$G and the pulsar age is $10^4 \\tau_4$y.  The broadband spectral energy distribution (SED, Figure 4) of our PWN candidate provides some useful constraints on its nature, if we interpret the radio/X-ray spectrum as synchrotron emission from a cooling electron population: ${\\rm d} N_e/{\\rm d}\\Gamma_e = A \\Gamma_e^{-s}$ with $\\Delta s= 1$ at a break $\\Gamma_B$. This corresponding to a photon break frequency $\\nu_B = 4.2 \\times 10^{13} B_{10} \\Gamma_6^2$, with a typical PWN field of $10B_{10} \\mu$G and a electron break at a Lorentz factor $\\Gamma_B = 10^6\\Gamma_6$. Our SED allows a modest range for the power law index $s=2.85\\pm 0.15$. The data require a very tight correlation between $s$ and the break frequency, ${\\rm Log} (\\nu_B) = 9.1\\, s -13.3$; however this allows a large range of break energies from the microwave band through the near IR.  The conventional PWN picture identifies this break with the cooling of the electrons at $\\tau_4 = 19 B_{10}^{-2} \\Gamma_6^{-1}$, associating this with the pulsar characteristic age. In this approximation (a homogenous, uniform field PWN) our data then can be used to eliminate $\\Gamma_B$ in favor of $s$, giving a PWN field estimate of $$ {\\rm Log}(B_{10}) \\approx 0.73 - 2/3 {\\rm Log}(\\tau_4) - 3.03 (s-3). $$ The total energy of the electron population, $\\approx 10^{47} B_{10}^{(s+1)/2} {\\rm erg}$, gives an upper limit on the initial pulsar period of $P_i \\la 0.4 B_{10}^{-(s+1)/4}$s. Of course, adiabatic losses in the PWN flow likely decrease the present electron energy and require smaller $P_i$.  \\begin{inlinefigure} \\epsscale{0.8} \\rotatebox{270}{\\plotone{f4.eps}} \\figcaption{Our PWN candidate's broad-band SED, with X-ray error regions for the full nebula (\\ASCA) and bright \\CXO\\ core.  The curve is the best-fit $s=2.85$ cooling synchrotron model.} \\label{SED} \\end{inlinefigure}  We can check consistency with another crude estimate of the nebula field.  \\citet{oka99} argued that the `fit' of the PWN candidate with Lynds 227 suggests confinement and pressure equilibrium. The CO line-width inferred pressure of $\\sim 1.4 \\times 10^{-11} {\\rm g/cm/s^2}$ implies a nebular equipartition field of $\\sim 20\\mu$G. However, since Lynds 227 brackets only on one side, the wind is likely only partly confined and $B$ may be larger.  For our best fit $s$, the equipartition $B$ suggests $\\tau_4 \\sim 20$, a Geminga-like pulsar. Our present poor constraint on $\\nu_B$ allows ages $\\sim 3 \\times 10^4 - 10^6$y, including $B_{12}\\tau_4 \\approx 10$ as estimated above. If the PWN $B$ exceeds $20\\mu$G, even smaller $\\tau$ are permitted.  We can also follow the pulsar surface field $B_{12}$ out to the termination shock in a simple 1-D picture. The magnetic field at $r_s = \\theta_s d$, the wind shock radius, is $B_s \\sim 3 B_\\ast r_\\ast^3/(r_{LC}^2 r_s ) \\sim 130 (B_{12} \\tau_4 d_{1.8} \\theta\\arcmin )^{-1}\\mu$G. Beyond $r_s$ the field evolution depends on the poorly understood wind magnetization, but assuming $B_s \\approx 20 \\mu$G constant in the postshock wind, for $B_{12}\\tau_4 \\approx 10$ we expect termination at $\\theta_s \\approx 40\\arcsec$. Our \\ACIS\\ image does not show a spherical wind with a subluminous zone at this $\\theta_s$, but instead suggests a jet or pulsar trail, with the $\\sim 15\\arcsec$ knot in the jet/trail plausibly a termination shock. The overall diffuse X-ray morphology suggests a pulsar trail leading back to a PWN, as seen for PSR B1757-24 \\citep{kasp01}, implying a birthsite to the NW. However, the nebula might also be powered through a jet, as for the asymmetric PSR 1509-58 PWN \\citep[c.f.][]{gae01}. In this case the faint extended emission to the SE (Figure 1) may be the counterlobe. Under this interpretation, the S32 cluster at $r \\approx 5 d_{1.8}$pc provides a plausible pulsar birthsite, requiring a pulsar transverse velocity $v_\\perp \\approx 500 d_{1.8}/\\tau_4$km/s. Much higher S/N imaging or a proper motion measurement are needed to distinguish these possibilities.  The $\\gamma$-ray flux is clearly an additional spectral component (Figure 4), but its origin is puzzling. 3EG J1809-2328 (= GeV J1809-2327) is apparently one of the more variable Galactic plane sources \\citep{tom99}. \\citet{oka99} suggested that the PWN-generated electrons penetrate the molecular cloud, generating GeV photons via relativistic Bremsstrahlung. As the PWN is well separated from the molecular gas at $\\sim (3\\arcmin) d \\sim 1.6$pc, one might expect variability no faster than the $3l/c \\sim 15$y flow-crossing time. This is substantially longer than the \\EGRET\\ variability timescale, suggesting an origin closer to the pulsar. Inverse Compton emission from the PWN termination shock seems attractive, but both Galactic IR and local synchrotron emission fluxes fail by several orders of magnitude to account for the required soft photon energy density. Even the optical/UV emission from the OB stars in S32 contribute less than $10^{5.5}L_\\odot$, failing to produce the required target photon density at the $\\sim 5$pc PWN distance by $\\ga 300 \\times$. The $\\gamma$-ray flux is quite plausible for magnetospheric (curvature) pulsar emission at the inferred ${\\dot E}$ and distance, but variability is unexpected from an isolated pulsar. Confirmation or exclusion of the $\\gamma$-ray variability, together with detection or limits on $\\gamma$-ray pulsations is the best way to address this issue. Unless an X-ray period can be measured, this will not be resolved until until AGILE or GLAST make new sensitive $\\gamma$-ray observations."
    },
    "0201/astro-ph0201184_arXiv.txt": {
        "abstract": "We present an atlas of spectra of O- and B-type stars, obtained with the Short Wavelength Spectrometer (SWS) during the Post-Helium program of the Infrared Space Observatory (ISO). This program is aimed at extending the Morgan \\& Keenan classification scheme into the near-infrared. Later type stars will be discussed in a seperate publication. The observations consist of 57 SWS Post-Helium spectra from 2.4 to 4.1 \\micron, supplemented with 10 spectra acquired during the nominal mission with a similar observational setting. For B-type stars, this sample provides ample spectral converage in terms of subtype and luminosity class. For O-type stars, the ISO sample is coarse and therefore is complemented with 8 UKIRT L'-band observations. In terms of the presence of diagnostic lines, the L'-band is likely the most promising of the near-infrared atmospheric windows for the study of the physical properties of B stars. Specifically, this wavelength interval contains the \\bra, \\pfg, and other Pfund lines which are probes of spectral type, luminosity class and  mass loss. Here, we present simple empirical methods based on the lines present in the 2.4 to 4.1 \\mum\\ interval that allow the determination of {\\em i)}   the spectral type of B dwarfs and giants to within two subtypes; {\\em ii)}  the luminosity class of B stars to within two classes; {\\em iii)} the mass-loss rate of O stars and B supergiants to within 0.25 dex. ",
        "introduction": "\\begin{table*}[!ht] \\begin{center} \\caption{The 12 O-type stars and one Wolf-Rayet star observed during the ISO/SWS Post-Helium mission, supplemented with 7 O stars observed with CGS4/UKIRT. The spectrum averaged signal-to-noise ratio (S/N) is listed in the last column.} \\label{OB1} \\begin{tabular}{lclcccr} \\hline \\hline Star&Name&Spectral&Spect. Type             &ISO Observation&Instrument&S/N\\\\ &    &Type    &Reference$^{\\mathrm{a}}$&   Number      &          &   \\\\ \\hline \\object{HD~46223  }&NGC 2244 203  &O4V((f))    &W72&            &UKIRT    &220\\\\ \\object{HD~190429A}&              &O4If+       &W73&89300401    &ISO/PHe  & 15\\\\ \\object{HD~46150  }&NGC 2244 122  &O5V(f)      &W72&            &UKIRT    &150\\\\ \\object{HD~199579 }&HR8023        &O6V((f))    &W73&89300301    &ISO/PHe  & 20\\\\ \\object{HD~206267 }&HR8281        &O6.5V((f))  &W73&90001601    &ISO/PHe  & 30\\\\ \\object{HD~47839  }&15 Mon        &O7V((f))    &W72&            &UKIRT    & 70\\\\ \\object{HD~24912  }&$\\zeta$ Per   &O7.5III((f))&W73&            &UKIRT    &125\\\\ \\object{HD~188001 }&QZ Sge        &O7.5Iaf     &W72&90000801    &ISO/PHe  & 10\\\\ \\object{HD~36861  }&$\\lambda$ Ori A&O8III((f)) &W72&            &UKIRT    &180\\\\ \\object{HD~209481 }&LZ Cep        &O9V         &W73&90001701    &ISO/PHe  & 25\\\\ \\object{HD~193322 }&HR7767        &O9V((n))    &W72&88201401    &ISO/PHe  & 15\\\\ \\object{HD~37043  }&$\\iota$ Ori   &O9III       &W72&            &UKIRT    &140\\\\ \\object{HD~207198 }&HR8327        &O9Ib-II     &W72&88502001    &ISO/PHe  & 25\\\\ \\object{HD~38666  }&$\\mu$ Col     &O9.5V       &W73&90701901    &ISO/PHe  & 10\\\\ \\object{HD~37468  }&$\\sigma$ Ori  &O9.5V       &C71&            &UKIRT    &165\\\\ \\object{HD~209975 }&19 Cep        &O9.5Ib      &W72&90001501    &ISO/PHe  & 30\\\\ \\object{HD~188209 }&HR7589        &O9.5Iab     &W72&88000501    &ISO/PHe  & 20\\\\ \\object{HD~30614  }&$\\alpha$ Cam  &O9.5Ia      &W72&88300601    &ISO/PHe  & 85\\\\ \\object{HD~195592 }&              &O9.7Ia      &W72&90001101    &ISO/PHe  & 45\\\\ \\object{WR~147    }&              &WN8h        &S96&88000701    &ISO/PHe  & 35\\\\ \\hline \\end{tabular} \\begin{list}{}{} \\item[$^{\\mathrm{a}}$] C71: Conti \\& Alschuler 1971; W72: Walborn 1972; W73: Walborn 1973; S96: Smith et al. 1996. \\end{list} \\end{center} \\end{table*} \\begin{table*}[!ht] \\begin{center} \\caption{The 30 B-type stars observed during the ISO/SWS Post-Helium mission, supplemented with 3 stars from the ISO/SWS nominal mission and 1 star observed with CGS4/UKIRT.} \\label{OB2} \\begin{tabular}{lclcccr} \\hline \\hline Star&Name&Spectral&Spect. Type             &ISO Observation&Instrument&S/N\\\\ &    &Type    &Reference$^{\\mathrm{a}}$&   Number      &          &   \\\\ \\hline \\object{HD~202214 }&HR8119        &B0V         &M55&90300701    &ISO/PHe  & 15\\\\ \\object{HD~93030  }&$\\theta$ Car  &B0Vp        &B62&25900905    &ISO/Nom  &115\\\\ \\object{HD~37128  }&$\\epsilon$ Ori&B0Ia        &W90&            &UKIRT    &115\\\\ \\object{HD~198781 }&HR7993        &B0.5V       &M55&88301201    &ISO/PHe  & 10\\\\ \\object{HD~207793 }&              &B0.5III     &M55&88700901    &ISO/PHe  & 20\\\\ \\object{HD~185859 }&HR7482        &B0.5Ia      &M55&89901301    &ISO/PHe  &  6\\\\ \\object{HD~116658 }&$\\alpha$ Vir  &B1V         &M55&25302001    &ISO/Nom  &165\\\\ \\object{HD~208218 }&              &B1III       &M55&88701101    &ISO/PHe  &  7\\\\ \\object{HD~190066 }&              &B1Iab       &M55&88101401    &ISO/PHe  & 15\\\\ \\object{HD~158926 }&$\\lambda$ Sco &B1.5IV      &H82&49101016    &ISO/Nom  &140\\\\ \\object{HD~52089  }&$\\epsilon$ Cma&B1.5II      &W90&88602001    &ISO/PHe  &130\\\\ \\object{HD~194279 }&V2118 Cyg     &B1.5Ia      &L92&88201301    &ISO/PHe  & 80\\\\ \\object{HD~193924 }&$\\alpha$ Pav  &B2IV        &L75&88500501    &ISO/PHe  & 95\\\\ \\object{HD~206165 }&9 Cep         &B2Ib        &L68&88300301    &ISO/PHe  & 70\\\\ \\object{HD~198478 }&55 Cyg        &B2.5Ia      &L68&88100501    &ISO/PHe  &100\\\\ \\object{HD~160762 }&$\\iota$ Her   &B3V         &J53&89900101    &ISO/PHe  & 70\\\\ \\object{HD~207330 }&$\\pi^{2}$ Cyg &B3III       &M55&88701301    &ISO/PHe  & 45\\\\ \\object{HD~15371  }&$\\kappa$ Eri  &B5IV        &H69&90701401    &ISO/PHe  & 35\\\\ \\object{HD~184930 }&$\\iota$ Aql   &B5III       &L68&88000901    &ISO/PHe  & 45\\\\ \\object{HD~191243 }&HR7699        &B5II        &L92&88401401    &ISO/PHe  & 20\\\\ \\object{HD~58350  }&$\\eta$ Cma    &B5Ia        &W90&90702301    &ISO/PHe  & 90\\\\ \\object{HIC~101364}&Cyg OB2 12    &B5Ia        &M91&90300901    &ISO/PHe  &105\\\\ \\object{HD~203245 }&HR8161        &B6V         &L68&88701401    &ISO/PHe  & 10\\\\ \\object{HD~155763 }&$\\zeta$ Dra   &B6III       &L68&89900201    &ISO/PHe  & 80\\\\ \\object{HD~209952 }&$\\alpha$ Gru  &B7IV        &H69&88500701    &ISO/PHe  &150\\\\ \\object{HD~183143 }&HT Sge        &B7Ia        &M55&89901501    &ISO/PHe  & 60\\\\ \\object{HD~14228  }&$\\phi$ Eri    &B8V-IV      &H69&88701901    &ISO/PHe  & 75\\\\ \\object{HD~207971 }&$\\gamma$ Gru  &B8III       &H82&88500901    &ISO/PHe  &100\\\\ \\object{HD~208501 }&13 Cep        &B8Ib        &L92&88701201    &ISO/PHe  & 80\\\\ \\object{HD~199478 }&V2140 Cyg     &B8Ia        &L92&88501801    &ISO/PHe  & 60\\\\ \\object{HD~16978  }&$\\epsilon$ Hyi&B9V         &H75&88401901    &ISO/PHe  & 65\\\\ \\object{HD~196867 }&$\\alpha$ Del  &B9IV        &M73&88101701    &ISO/PHe  & 75\\\\ \\object{HD~176437 }&$\\gamma$ Lyr  &B9III       &J53&88401501    &ISO/PHe  &110\\\\ \\object{HD~202850 }&$\\sigma$ Cyg  &B9Iab       &M55&90600601    &ISO/PHe  & 65\\\\ \\hline \\end{tabular} \\begin{list}{}{} \\item[$^{\\mathrm{a}}$] J53: Johnson \\& Morgan 1953. M55: Morgan et al. 1955; B62: Buscombe 1962; L68: Lesh 1968; H69: Hiltner et al. 1969; M73: Morgan \\& Keenan 1973; L75: Levato 1975; H75: Houk 1975; H82: Houk 1982; W90: Walborn \\& Fitzpatrick 1990; M91: Massey \\& Thompson 1991; L92: Lennon et al. 1992. \\end{list} \\end{center} \\end{table*}  \\begin{table*}[!ht] \\begin{center} \\caption{The 14 B stars with emission lines observed in the ISO/SWS Post-Helium mission supplemented with 7 stars observed during the ISO/SWS nominal mission.} \\label{Be} \\begin{tabular}{lclcccr} \\hline \\hline Star&Name&Spectral&Spect Type              &Observation &Status&S/N\\\\ &    &Type    &Reference$^{\\mathrm{a}}$&   Number   &      &   \\\\ \\hline \\object{V1478~Cyg }&MWC 349A      &O9III[e]  &Z98&18500704    &ISO/Nom  &145\\\\ \\object{HD~206773 }&MWC 376       &B0Vpe     &M55&88502101    &ISO/PHe  & 20\\\\ \\object{HD~5394   }&$\\gamma$ Cas  &B0.5Ve    &P93&24801102    &ISO/Nom  &150\\\\ \\object{HD~212571 }&$\\pi$ Aqr     &B1Ve      &L68&90601301    &ISO/PHe  & 20\\\\ \\object{HD~50013  }&$\\kappa$ Cma  &B1.5IVne  &H69&90702001    &ISO/PHe  & 80\\\\ \\object{HD~200775 }&MWC 361       &B2V[e]    &G68&90300501    &ISO/PHe  & 50\\\\ \\object{HD~45677  }&MWC 142       &B2V[e]    &Z98&71101992    &ISO/Nom  &135\\\\ \\object{HD~56139  }&$\\omega$ Cma  &B2IV-Ve   &H69&90702201    &ISO/PHe  & 40\\\\ \\object{HD~105435 }&HR 4621       &B2IVne    &H69&07200272    &ISO/Nom  &120\\\\ \\object{HD~205021 }&$\\beta$ Cep   &B2IIIe    &M55&88100301    &ISO/PHe  &110\\\\ \\object{HD~187811 }&12 Vul        &B2.5Ve    &L68&90700901    &ISO/PHe  & 25\\\\ \\object{HD~191610 }&28 Cyg        &B2.5Ve    &L68&89900901    &ISO/PHe  & 35\\\\ \\object{HD~205637 }&$\\epsilon$ Cap&B3Vpe     &H88 &90601701    &ISO/PHe  & 30\\\\ \\object{HD~10144  }&$\\alpha$ Eri  &B3Vpe     &H69&90000101    &ISO/PHe  &140\\\\ \\object{HD~56014  }&EW Cma        &B3IIIe    &H82 &90702101    &ISO/PHe  & 20\\\\ \\object{HD~50123  }&HZ Cma        &B6Vnpe    &S   &88601901    &ISO/PHe  & 75\\\\ \\object{HD~198183 }&$\\lambda$ Cyg &B6IVe     &L68&89900801    &ISO/PHe  & 35\\\\ \\object{HD~209409 }&omi Aqr       &B7IVe     &L68&90601501    &ISO/PHe  & 35\\\\ \\object{HD~193237 }&P Cygni       &B2pe      &L68&33504020    &ISO/Nom  &100\\\\ \\object{HD~94910  }&AG Car        &B2pe      &H75&22400153    &ISO/Nom  & 80\\\\ \\object{HD~93308  }&$\\eta$ Car    &Bpe       &H75&07100250    &ISO/Nom  &170\\\\ \\hline \\end{tabular} \\begin{list}{}{} \\item[$^{\\mathrm{a}}$] M55: Morgan et al. 1955; L68: Lesh 1968; G68: Guetter 1968; H69: Hiltner et al. 1969; H75: Houk 1975; H82: Houk 1982; H88: Houk 1988; P93: van Paradijs 1993; Z98: Zorec et al. 1998; S: Simbad. \\end{list} \\end{center} \\end{table*}  The advance of infrared-detector technology since the eighties has opened new perspectives for the study of early-type stars. Investigation of the early phases of their evolution especially benefits from infrared (IR) observations. The birth places of massive stars are identified with Ultra-Compact H{\\sc ii} regions ({\\sc UCHii}). In such regions, the stars are still embedded in material left over from the star formation process and are obscured at optical and ultraviolet wavelengths. In the K-band (ranging from 2.0 to 2.4 \\mum) dust optical depths $\\tau$ of a few occur, while in the H-band (ranging from 1.5 to 1.8 \\mum) $\\tau$ is typically of order ten. At shorter wavelengths, the dust extinction becomes too high to observe the embedded stars. The IR emission of the warm dust cocoon covering the newly formed massive stars in {\\sc UCHii} regions peaks typically at about 100 \\mum. At wavelengths longwards of 5--10 \\mum, the thermal emission of the dust dominates the photospheric flux, and can be as much as 4 orders of magnitude above the stellar free-free continuum at 100 \\mum\\ (Churchwell 1991).  Reliable values for the luminosities, temperatures and mass-loss rates of the embedded massive stars are essential as they allow us to trace the very early phases of their evolution of which little is known. Furthermore, these parameters control the photo-dissociation and ionisation of the molecular gas, the evaporation of the dust, and affect the morphology of the {\\sc UCHii} region.  The development of {\\em quantitative} diagnostics based on IR spectral data requires, as a first step, homogeneous observations of a large set of both normal and peculiar non-embedded early-type stars, that have been studied in detail at optical and ultraviolet wavelengths where OB-type stars exhibit many spectral lines. Such stars may be used to calibrate quantitative methods based on IR spectroscopy alone. Calibration work has already been carried out in other near-infrared wavelength ranges, in the J-band e.g. Wallace et al. (2000), in the H-band e.g. Meyer et al. (1998) and Hanson et al. (1998), and in the K-band Hanson et al. (1996). The ``Post-Helium program'' conducted with the Short Wavelength Spectrometer (SWS) on board the Infrared Space Observatory (ISO) is intended to provide such a data set. This mission started after helium boil-off in April 1998 and made use of the ability of the detectors of SWS to acquire observations in band~1 [2.4-4.1] \\mum\\ during the slow warming of the satellite (see also Sect.~\\ref{ISO}) The band~1 of ISO SWS ranges from 2.4 to 4.1 \\mum, and is, like the K-band, positioned favourably in the narrow window in which newly born stars can be observed directly. This wavelength region contains important diagnostic hydrogen lines of the Brackett (\\bra, \\brb), Pfund (\\pfg), and Humphreys series.  In this paper, we present and study 75 spectra of early-type stars, 67 [2.4-4.1 \\mum] ISO/SWS spectra and 8 [3.5-4.1 \\mum] spectra observed with the United Kingdom Infrared Telescope (UKIRT). This sample includes OB, Be, and Luminous Blue Variable (LBV) stars. We discuss line trends as a function of spectral type, following a strategy similar to the one adopted by Hanson et al. (1996) for the K-band. Simple empirical methods are employed to derive the spectral type and/or luminosity class. These methods may also be applied if only ground-based L'-band spectra are available (which cover a smaller wavelength range).  The paper is organised as follows: In Sect.~\\ref{observation} we discuss the data acquisition and reduction techniques; Sect.~\\ref{atlas} comprises a catalogue of good quality spectra; Sect.~\\ref{identification} provides the line identifications. Line trends and methods to classify OB-type stars are presented in Sect.~\\ref{trends}, while Sect.~\\ref{emmission} describes the spectra of B stars with emission lines. The results are summarised in the final section. The equivalent-width measurements are listed in the Appendix. ",
        "conclusions": "\\label{summary} In this paper, we have presented an atlas of 2.4 to 4.1 \\mum\\ ISO SWS spectra of early-type stars, mainly obtained during its Post-Helium mission, and several 3.5-4.1 \\mum\\ spectra of O stars obtained at UKIRT. The observations include normal OB, Be and Luminous Blue Variable stars. Later spectral types will be presented in a separate publication (Vandenbussche et al. in prep.). We have explored a number of simple empirical methods aimed at using the infrared spectrum to {\\em i)} determine the spectral type and/or luminosity class, and {\\em ii)} determine the mass-loss rate. The main results are:  \\begin{enumerate}  \\item In normal B-type giant to dwarf stars the Pfund lines, and to a lesser extent the Brackett lines, may be used to estimate the spectral type. We provide a simple formula to do this. The leading line of each series shows the most pronounced dependence. Helium lines help to improve this spectral classification, \\hei\\ being present in late O-type and early B-type stars. All B-type giants and dwarfs have \\bra\\ in absorption. The full width at half maximum of this line may be used to discriminate between luminosity classes III and V, the line being broader for dwarfs. \\\\  \\item  In B-type supergiants the equivalent width of all measured hydrogen lines remain constant with spectral sub-type, although with a significant scatter. \\bra\\ is seen mostly in emission, while all other lines are in absorption. \\hei~$\\lambda$3.0736 is systematically stronger in absorption compared to B-type dwarfs and giants.\\\\  \\item In normal O-type stars and in B-type supergiants, the \\bra\\ line is mostly in emission and provides a sensitive indicator of the mass-loss rate. We give a relation that uses the equivalent width of this line to estimate $\\dot{M}$.\\\\   \\item Concerning hydrogen lines, the ones positioned in the L'-band seem best suited to derive physical properties of OB stars when compared to the diagnostics available in other atmospheric bands such as K-, H-, and J-band. The main reason is that the L'-band contains three different hydrogen series lines and includes the leading Brackett-series line. Concerning other species, the K-band seems to contain the most useful lines. This last remark, however, only applies to O-type stars (where e.g. C~{\\sc iv}, N~{\\sc iii} and an unblended He~{\\sc ii} line are seen) and not to B-type stars which do not show lines of metal species in that wavelength range.\\\\  \\item  In our sample of Be, B[e], and Luminous Blue Variable stars we find no obvious correlation between spectral type and strength of the emission lines. Stars with spectral type earlier than B3 show \\hei\\ lines, similar to normal B-type stars. Several emission line stars show \\oi, however not at spectral types later than B2.\\\\  \\end{enumerate}"
    },
    "0201/astro-ph0201467_arXiv.txt": {
        "abstract": "We present a method based on two broad-band colors to investigate the nature of the Extremely Red Objects (EROs), i.e., the galaxies selected to have very red optical-to-infrared colors. Dusty starburst and old ellipticals at redshifts between 1 and 2 appear to occupy two different regions of the J-K vs. I-K color diagram, allowing for an easy classification. This diagnostic was applied to a complete sample of 57 EROs: the two populations are found to be present in the sample in similar abundances. The cosmic star formation density in the dusty starbursts is found to be of the order of that in the Lyman-Break Galaxies (LBG). ",
        "introduction": " ",
        "conclusions": ""
    },
    "0201/astro-ph0201099_arXiv.txt": {
        "abstract": "A bound on the compactness of the neutron star in the low mass x-ray binary 4U 1636-53 is used to estimate the equation of state of neutron star matter at high density. Observations of 580 Hz oscillations during the rising phase of x-ray bursts from this system appear to be due to two antipodal hot spots on the surface of an accreting neutron star rotating at 290 Hz, implying the compactness of the neutron star is less than 0.163 at the 90\\% confidence level. The equation of state of high density neutron star matter estimated from this compactness limit is significantly stiffer than extrapolations to high density of equations of state determined by fits of experimental nucleon-nucleon scattering data and properties of light nuclei to two- and three-body interaction potentials. ",
        "introduction": "Oscillations in x-ray brightness during bursts have been observed from ten low mass x-ray binaries with the \\textit{Rossi} X-ray Timing Explorer \\cite{Strohmayer2001}. These oscillations appear to be due to rotational modulation of one or a pair of antipodal hotspots on the surface of the accreting neutron star \\cite{Strohmayer-etal-1997}. With this model the amplitude of the oscillations is determined by the mass-to-radius ratio of the neutron star, the compactness $M/R$. A more compact neutron star results in greater deflection of the x-ray photons emanating from a hot spot, making the hot spot visible to an observer for a larger portion of the rotation period and so reducing the amplitude of the brightness oscillations. The bounds on the compactness of the neutron star are particularly limiting when two antipodal hot spots are present \\cite{MillerAndLamb1998}.  The source 4U 1636-53 appears to be the best candidate for providing a stringent limit on the compactness of neutron stars from x-ray burst oscillations \\cite{Nath-etal-2001}. A 290 Hz subharmonic of the stronger 580 Hz brightness oscillation has been observed in five bursts from this source \\cite{Miller1999}, but the subharmonic has not been confirmed in subsequent bursts \\cite{Strohmayer2001}. The existence of the subharmonic would suggest the spin frequency of the neutron star is 290 Hz, and the presence of two antipodal hot spots is responsible for the 580 Hz modulation. A pair of high frequency quasiperiodic oscillations separated by 251 Hz has also been reported for this source and interpreted in a beat frequency model to be a consequence of a spin frequency of 290 Hz with two hot spots on the neutron star \\cite{Miller-etal-1998}. The x-ray oscillations during the rising phase of bursts from 4U 1636-53 have been recently modeled as due to a pair of circular antipodal hot spots, each having an angular size that grows linearly in time \\cite{Nath-etal-2001}. The two hot spot model constrains the neutron star compactness to $M/R < 0.163$ at 90\\% confidence (using $G = c = 1$), requiring the radius of a 1.4 $M_\\odot$ neutron star to be greater than 12.7 km. The 99\\% confidence level compactness limit, $M/R < 0.183$, implies $R > 11.3$ km when $M = 1.4 M_\\odot$.   During a burst event in low mass x-ray binaries the oscillation frequency is usually observed to evolve asymptotically to a limit that probably corresponds to the spin rate of the neutron star. Measurement of the asymptotic frequency at different orbital phases of the binary could reveal a Doppler modulation that would allow for a determination of the masses of the neutron star and the companion. A mass measurement together with the compactness limit would tightly constrain the radius of the star. A Doppler modulation has not been observed for 4U 1636-53 in an analysis of 26 burst oscillations \\cite{Giles-etal-2001}. The absence of the modulation likely limits the neutron star mass in this system to $M < 1.6 M_\\odot$ \\cite{Giles-etal-2001}.  The purpose of this paper is to use the 90\\% confidence level compactness bound on the 4U 1636-53 neutron star to estimate the high density equation of state of neutron star matter. Ideally one would have both a measured mass and radius for a number of neutron stars over a wide range of masses up to the maximum mass value. The observed mass-radius relationship could then be used to uniquely calculate the equation of state by inverting the Oppenhemier-Volkoff equation, the relativistic stellar structure solution to Einstein's equations \\cite{Lindblom1992}. Even if the mass and radius are known for only a single neutron star, the high density equation of state can still be accurately estimated up to the central density yielding the mass of this star. At low density the equation of state can be determined by fitting experimental nucleon-nucleon scattering data and properties of light nuclei to two-body and three-body interaction potentials \\cite{Wiringa-etal-1988, Akmal-etal-1998}. Confidence in these low density equations of state is high near normal nuclear density (energy density = 153 MeV/fm$^3$, baryon density = 0.16 fm$^{-3}$) because nuclear physics experiments allow for verification in the low density regime. Estimates of the maximum possible mass of neutron stars based on these equations of state have regarded them as experimentally verified up to twice normal nuclear density \\cite{KalogeraAndBaym1996, Olson2001}. Assuming one has values of both the mass and radius for a single neutron star, the equation of state in the density range between the experimentally verified low density regime and the high central density of this star can be accurately estimated to have a polytrope form provided the matter does not undergo a phase transition in this density range \\cite{Lindblom1992}. The polytrope high density equation of state is chosen to match the energy density and pressure of the low density equation of state at the upper limit of applicability of the low density equation, and the adiabatic index of the polytrope is chosen to yield a stellar radius corresponding to the observed radius of the star when the mass of the model star equals the observed mass. The resulting polytrope is an estimate of the equation of state of neutron star matter in the high density regime.  In the case of the neutron star in 4U 1636-53 we do not have measurements of both its mass and radius, but only an upper bound on the compactness and an upper limit on the mass. Nevertheless the high density equation of state is still tightly constrained by the compactness limit provided it is assumed the neutron star mass is at least $1.4 M_\\odot$, a plausible assumption for an old accreting system like 4U 1636-53. In the next section the high density equation of state is estimated from the 90\\% confidence level compactness bound on the neutron star in 4U 1636-53 assuming the neutron star mass is at least $1.4 M_\\odot$. The main result of this paper is that the A18 + $\\delta v$ + UIX* equation of state of Ref. \\cite{Akmal-etal-1998} (hereafter the APR equation of state), a recently developed equation of state representative of the class of equations of state derived from fits of experimental nuclear physics data to interaction potentials, is too soft at high density to yield a star with the observed low compactness of the neutron star in 4U 1636-53. Further if the APR equation of state is used up to twice normal nuclear energy density, and any polytrope equation of state having a sound speed less than the speed of light is used above twice normal nuclear density, the model star compactness still exceeds the compactness of the 4U 1636-53 neutron star. Using the APR equation of state up to just a value in the range 1.55-1.64 times normal nuclear energy density and a stiff polytrope with an adiabatic index $\\Gamma$ of 5.5 $\\alt \\Gamma \\alt$ 7 at high density can produce a star with the low compactness of 4U 1636-53 without the speed of sound waves in the neutron star interior exceeding the speed of light. ",
        "conclusions": "A new equation of state (Table \\ref{EOS}) valid for describing neutron star matter up to the central density of a 1.4 $M_\\odot$ star has been developed in this paper based on observations of the low mass x-ray binary 4U 1636-53. The validity of this equation of state rests on three main assumptions: the actual compactness of the 4U 1636-53 neutron star does not exceed the 90\\% confidence level bound, the mass of the 4U 1636-53 neutron star is near 1.4 $M_\\odot$, and the APR equation of state is valid up to near 1.64 times normal nuclear energy density. The APR equation of state used at all densities is sufficient to describe the structure of the 4U 1636-53 neutron star if its compactness is actually near the 99\\% confidence level bound, $M/R = 0.183$, rather than near the 90\\% confidence level bound value of $M/R = 0.163$. The APR equation of state used at all densities is also adequate to meet the 90\\% confidence level bound if the 4U 1636-53 neutron star is 1.27$M_\\odot$ or less. (The APR equation of state is seen from Fig. \\ref{Fig_1} to yield a radius of 11.5 km for a mass of 1.27$M_\\odot$, making the compactness of this star $M/R = 0.163$.) If the mass of the 4U 1636-53 neutron star is larger than 1.4$M_\\odot$, the equation of state of neutron star matter must be even stiffer than the Table \\ref{EOS} equation of state. The APR equation of state is based on a large body of experimental evidence at low density, so it is reasonable to trust its validity up to $\\tilde{\\rho} \\approx 1.6\\rho_{nn}$. Figure \\ref{Fig_3} shows the equation of state derived here (Table \\ref{EOS}) is relatively insensitive to the exact transition energy density value as long as $\\tilde{\\rho}$ is near 1.6$\\rho_{nn}$.  If the APR equation of state is used up to $1.55\\rho_{nn} \\alt \\tilde{\\rho} \\alt 1.64\\rho_{nn}$, a stiff polytrope, having an adiabatic index of 5.5 $\\alt \\Gamma \\alt$ 7, is needed to produce a model star with a compactness $M/R = 0.163$ without the speed of sound waves in the neutron star interior exceeding the speed of light. The adiabatic index would be small, near 5/3 for nonrelativistic nucleons and 4/3 in the relativistic limit, if the pressure is due mainly to the Fermi kinetic energy. The adiabatic index is near 2 if the equation of state is dominated by static two-body interactions, near 3 if static three-body interactions are most important, and in excess of 3 if repulsive momentum-dependent interactions are dominant \\cite{Akmal-etal-1998}. The large adiabatic index of the equation of state at high density implied by the 4U 1636-53 neutron star 90\\% confidence level compactness bound suggests the APR equation of state underestimates at high density the strength of the repulsive momentum-dependent interactions."
    },
    "0201/astro-ph0201050_arXiv.txt": {
        "abstract": "The motivations and the status of the K20 survey are presented. The first results on the evolution of massive galaxies and the comparison with the predictions of the currently competing scenarios of galaxy formation are also discussed. ",
        "introduction": "Understanding the evolution of massive galaxies (e.g. M$_{stellar}$\\gsim$10^{11}$ M$_{\\odot}$) is important because to constrain the different scenarios of structure and galaxy formation. In particular, the question on the formation of the present-day massive spheroidals is still one of the most debated issues of galaxy evolution (see \\cite{renz} for a review). In one scenario, massive spheroidals are formed at early cosmological epochs (e.g. $z>3$) through a short and intense episode of star formation (with $SFR\\sim 100-1000$ M$_{\\odot}$yr$^{-1}$), followed by a passive evolution (or pure luminosity evolution, PLE) of the stellar population to nowadays. In marked contrast, the hierarchical scenarios predict that massive spheroidals are the product of rather recent merging of pre-existing disk galaxies taking place mostly at lower redshifts and with moderate star formation rates \\cite{k96,bau}. In hierarchical merging scenarios, fully assembled massive field spheroidals with M$_{stellar}$\\gsim$10^{11}$ M$_{\\odot}$ at $z$\\gsim 1 are very rare objects \\cite {k98} (see also Baugh, this volume).  From an observational point of view, a direct way to test the above scenarios is to study the evolution of massive galaxies by means of spectroscopic surveys of field galaxies selected in the $K$-band \\cite{bro,cowie,cohen,stern}. Since the near-IR light is a good tracer of the galaxy mass\\cite{gava,k98}, $K$-band imaging allows to select massive galaxies at high-$z$. A galaxy with a stellar mass of about $10^{11}$ M$_{\\odot}$ is expected to have $18<K<20$ for $1<z<2$ \\cite{k98}, thus implying that moderately deep $K$-band surveys can efficiently select massive galaxies in that redshift range. Deep spectroscopy with 8-10m class telescopes can then be used to search for massive systems and to constrain their redshift distribution. ",
        "conclusions": ""
    },
    "0201/astro-ph0201266_arXiv.txt": {
        "abstract": "Observations with the Space Telescope Imaging Spectrograph aboard the Hubble Space Telescope have been used to constrain the atomic sulfur column density in Io's atmosphere.  The \\ion{S}{1} $\\lambda$1479 dipole allowed and forbidden transition multiplets have been resolved for the first time at Io, enabling the study of both optically thick and thin transitions from a single atomic species.  The allowed transitions contribute 62 $\\pm$ 8\\% and the forbidden transitions 38 $\\pm$ 8\\%, on average, to the total signal of the \\ion{S}{1} $\\lambda$1479 multiplets.  Using the optically thick and thin transitions of \\ion{S}{1} $\\lambda$1479 observed near the limbs of Io, we derive a tangential atmospheric sulfur column abundance of 3.6$\\times$10$^{12}$\\,cm$^{-2}\\,<\\,{\\cal N}_{s}\\,<$\\,\\,1.7$\\times$10$^{13}$\\,cm$^{-2}$, which is independent of electron temperature and density.  A low density SO$_2$ atmosphere, ${\\cal N}_{so_2}\\sim$\\,\\,5-10$\\times$10$^{15}$\\,cm$^{-2}$, consistent with that inferred from other recent observations, is most consistent with these bounds. ",
        "introduction": "Io's atmosphere plays a key role in populating the jovian magnetosphere with plasma and replenishing the plasma torus with sulfur and oxygen \\citep{spe96}.  This tenuous atmosphere, first detected by the Voyager 1 IRIS experiment over 20 years ago \\citep{pea79}, has proven to be a challenge to characterize quantitatively.  It is composed primarily of SO$_2$ \\citep{lel90,lel92} and SO \\citep{lel96}.  The atomic by-products of SO$_2$, sulfur and oxygen, were first identified in the Io plasma torus \\citep{bro81,dur83}, and more recently detected in Io's emission line spectrum \\citep{bal87}.  Sodium \\citep{bur98,gei99,bou00} and potassium \\citep{bro01} are also present at much lower abundance than sulfur and oxygen.  The observed emissions are produced when the atmospheric sulfur and oxygen are collisionally excited by impinging electrons from the plasma torus.  The emissions are therefore diagnostic of the plasma interaction with Io's atmosphere, providing quantitative information about the electron temperature and density, in addition to the abundances of sulfur and oxygen.  Despite the recent advances in quantifying Io's atmosphere, including investigating the inhomogeneity and impact of volcanic activity \\citep{fel00,mcg00,spe00}, analyzing the airglow \\citep{gei99}, and resolving the morphology \\citep{roe99,ret00}, and the fact that several sets of sulfur and oxygen emission line data have been acquired and analyzed \\citep{bal89,mcg00,oli01,wol01}, the abundances are still not well constrained.  Various observational measurements and theoretical models have placed loose bounds on the sulfur column density (${\\cal N}_{s}$). Previously, electron densities and temperatures were assumed in order to estimate ${\\cal N}_{s}$, resulting in limits of 2.2$\\times$10$^{12}$\\,cm$^{-2}\\,<\\,{\\cal N}_{s}\\,<\\,\\,$7$\\times$10$^{15}$\\,cm$^{-2}$ \\citep{bal87,mcg00}.  Recent observations make both optically thick and optically thin regimes available in a single exposure, allowing us to constrain ${\\cal N}_{s}$ without assuming a specific $T_e$ or $n_e$.  We present such a temperature and density independent method for determining the column density in this paper.  Sulfur and oxygen abundances also provide a little-utilized diagnostic of the origin and nature of the SO$_2$ atmosphere.  Significant progress has been made in characterizing both the spatial distribution and temporal variability of the SO$_2$ atmosphere \\citep{bal94,tra96,mcg00,fel00,spe00,str01}.  The current picture includes a higher column density equatorial, sub-solar region, ${\\cal N}_{so_2}\\sim\\,\\,$1-5$\\times$10$^{16}$\\,cm$^{-2}$, with a strong poleward gradient, where ${\\cal N}_{so_2}$ decreases to $<\\,$10$^{15}$\\,cm$^{-2}$ \\citep{fel00,str01} and column densities $>\\,$10$^{16}$\\,cm$^{-2}$ above active volcanic regions \\citep{mcg00,spe00}. It is obvious that both active volcanism and SO$_2$ frost sublimation are important atmospheric sources, although the relative importance of each mechanism has not been determined \\citep{sum96,won00,mos01}.  The models show that sulfur can serve as a proxy for SO$_2$, with enhanced abundance relative to oxygen during and shortly after volcanic activity as compared to a sublimation atmosphere.  Sulfur displays similar trends as, and scales well with, the SO$_2$ abundance in most of the dayside models.  Determining and monitoring the sulfur abundance, can, in principle, provide a much-needed measure of the variability of Io's atmosphere.  We describe the observations and explain the data reduction in Section 2.  We then give the motivation for and present the analysis of the opacity of the sulfur transitions and the atomic sulfur column density in Section 3.  Finally, we discuss the implications of our findings in Section 4. ",
        "conclusions": "The STIS observations presented here and used to calculate the atomic sulfur column density are tangential cuts through Io's atmosphere centered on the equatorial spots; therefore, the column densities are also tangential. \\citet{wol01} has recently published a collection of four years of HST/STIS low resolution observations of near Io emissions and extended emissions in the 1150-1730 \\AA\\, wavelength range.  The photon fluxes measured at 1 Io radius for the \\ion{S}{1} $\\lambda$1479 multiplets varied between 3$\\times$10$^{-4}$\\,cm$^{-2}$\\,sec$^{-1}$ and 3$\\times$10$^{-3}$\\,cm$^{-2}$\\,sec$^{-1}$, corresponding to eclipse and western elongation, respectively.  These values are consistent with our HST/STIS photon fluxes for \\ion{S}{1} $\\lambda$1479, implying that the tangential sulfur column that \\citet{wol01} were sampling should have values for ${\\cal N}_{s}$ similar to ours.  \\citet{mcg00} determined the vertical column density of S, SO, and SO$_2$ above specific locations on Io, including the Pele volcano.  With an estimated collision strength and electron impact excitation rate coefficient for the \\ion{S}{1}] $\\lambda$1900 doublet, and assuming canonical plasma torus values of electron density and temperature, 1000 cm$^{-3}$ and 5 eV, they determine an ${\\cal N}_{s}$ of 1$\\times$10$^{14}$\\,cm$^{-2}$ at Pele, which is much larger than our upper bound of 1.3$\\times$10$^{13}$\\,cm$^{-2}$.  A column density that is significantly larger than our value should result in additional optical depth effects.  In Figure 6, the Pele flux spectrum published in Figure 2 of \\citet{mcg00} is compared with the optically thin FOS synthetic spectrum described earlier, and confirms that the \\ion{S}{1} $\\lambda$1814 multiplet is more optically thick over Pele than in our IUE and FOS spectra, lending credence to the significantly higher value of ${\\cal N}_{s}$ over Pele despite the uncertainty in the \\ion{S}{1}] $\\lambda$1900 electron excitation rate. This implies that ${\\cal N}_{s}$ is larger in volcanic plumes than in the equatorial spots.  In the 1-D, steady state photochemical atmospheric models presented by \\citet{sum96}, both high and low density SO$_2$ atmosphere sub-solar profiles with various levels of vertical eddy mixing are calculated.  We have extracted the vertical column densities of several atmospheric components from their published profiles.  The ${\\cal N}_{s}$ values corresponding to the high density SO$_2$ cases, ${\\cal N}_{so_2}$ = 10$^{18}\\,$cm$^{-2}$, were much higher than the limits from our analysis.  For example, with a low vertical eddy mixing coefficient of $k$ = 10$^{6}\\,$cm$^{2}$\\,sec$^{-1}$, ${\\cal N}_{s}$ $\\sim$ 2$\\times$10$^{17}\\,$cm$^{-2}$; for a high vertical eddy mixing coefficient of $k$ = 10$^{9}\\,$cm$^{2}$\\,sec$^{-1}$, ${\\cal N}_{s} \\sim\\,\\,$9$\\times$10$^{14}\\,$cm$^{-2}$.  On the other hand, for their low density case (${\\cal N}_{so_2}$ = 8$\\times$10$^{15}\\,$cm$^{-2}$ with high eddy mixing $k$ = 10$^{9}\\,$cm$^{2}$\\,sec$^{-1}$) ${\\cal N}_{s}$ $\\sim$ 6$\\times$10$^{12}\\,$cm$^{-2}$.  The corresponding tangential column density, assuming spherical symmetry and using the atmospheric sulfur number density profile in \\citet{sum96}, is on the order of 4$\\times$10$^{13}\\,$cm$^{-2}$. This value is much more consistent with our results and the general overall picture of the SO$_2$ atmosphere described in the introduction.  In a 3-D high density SO$_2$ model atmosphere presented by \\citet{won00}, they assume an initial sub-solar ${\\cal N}_{so_2}$ of 6$\\times\\,$10$^{17}\\,$cm$^{-2}$, which results in a tangential sulfur column density at the terminator, ${\\cal N}_{s}\\sim\\,\\,$2$\\times$10$^{15}\\,$cm$^{-2}$, inconsistent with our upper limit to ${\\cal N}_{s}$.  In the model in which the input SO$_2$ column is reduced to a lower density value of 7$\\times\\,$10$^{16}\\,$cm$^{-2}$, the tangential ${\\cal N}_{s}$ is more consistent with our results at a value of 6$\\times\\,$10$^{13}\\,$cm$^{-2}$.  In order to predict what the atmospheric driving mechanism is, \\citet{mos01} have modeled both an SO$_2$ frost sublimation atmosphere and a Pele-type volcanically driven atmosphere.  Of the two models, our ${\\cal N}_{s}$ value is more consistent with the SO$_2$ frost sublimation atmosphere simply because this model has a lower sulfur column density. However, the true test to the Moses et al. (2001) models is to compare the sulfur to oxygen column abundances extracted from a single exposure.  When the sulfur is more abundant than the oxygen, the models imply a volcanically driven atmosphere.  None of the medium resolution data presented here had a simultaneous oxygen detection, so the ratio is not currently available, although this sulfur to oxygen relationship will be analyzed in future work with a more extensive data set."
    },
    "0201/astro-ph0201500_arXiv.txt": {
        "abstract": "Focussing on preplanetary grains growth, we discuss the properties of dust aggregation driven by magnetic dipole forces. While there is no direct evidence for the existence of magnetic grains present in the solar nebula, there are reasons to assume they may have been present.  We derive analytical expressions for the cross-section of two interacting dipoles.  The effective cross section depends upon the strength of the magnetic dipoles and the initial velocities.  For typical conditions the magnetic cross section is between 2 and 3 orders of magnitude larger than the geometric cross section.  We study the growth dynamics of magnetic grains and find that the mass of the aggregates should increase with time like t$^{3.2}$ whereas Brownian motion growth behaves like t$^2$. A numerical tool is introduced which can be used to model dust aggregation in great detail, including the treatment of contact forces, aggregate restructuring processes and long-range forces.  This tool is used to simulate collisions between magnetic grains or clusters and to validate the analytical cross-sections. The numerically derived cross section is in excellent agreement with the analytical expression.  The numerical tool is also used to demonstrate that structural changes in the aggregates during collisions can be significant. ",
        "introduction": "Coagulation of dust grains is generally believed to be the mechanism by which growth of particles in the early solar nebula proceeds from sub-micron grains all the way up to planetesimals \\citep{weid-cuzzi-PPIII,2000prpl.conf..533B}.  However, many questions remain about the detailed way in which coagulation proceeds, in particular during the initial phase when the particles are still very small and dynamically well coupled to the gas.  If the particles are very strongly coupled to the gas, collisions have to rely on the relative velocities induced by Brownian motion \\citep{kempf99:_n_partic_simul_dust_growt}.  While these slow collisions are certain to result in almost perfect sticking between the grains \\citep{Chokshi-ea93,poppe00:_aenal_exper_stick}, it turns out that they may be too slow  to produce the necessary dust growth in the limited time frame available \\citep{kempf99:_n_partic_simul_dust_growt}.  First of all, the low relative velocities severely limit the number of collisions taking place.  As a second complication, the energies involved in the collisions are insufficient to change the structure of the aggregates \\citep{coagu,blum00:_exper_stick_restr}.  Simple hit-and-stick growth in a cluster-cluster aggregation process leads to extremely fluffy aggregates with low fractal dimensions \\citep{me91,kempf99:_n_partic_simul_dust_growt}. \\citeauthor{kempf99:_n_partic_simul_dust_growt} showed that the Brownian motion leads to fractal dimensions $D<2$ for which significant changes in the cross-section-to-mass ratio can only be expected for very large particles.  The friction time, i.e. the time in which a particle adapts its velocity to the surrounding gas, changes in a Brownian growth process only as $r^{1/5}$ with the particle radius.  Furthermore, \\citeauthor{kempf99:_n_partic_simul_dust_growt} neglected the rotation of the aggregating particles which will lead to even lower fractal dimensions of the collision product \\citep{bl00}.  Since the aerodynamic behavior of particles is governed by the friction time \\citep{weid_aerodyn}, the aerodynamic properties of large fluffy aggregates and of small grains are almost identical.  This poses a serious problem for the coagulation process which must accelerate eventually, in order to fulfill the time constraints given.  Possible ways to accelerate growth have been discussed in the literature, in particular grain settling and radial drift \\citep[e.g.][]{Weid_settling}, or differential coupling to turbulent eddys of different sizes \\citep{Weidenschilling84,Mizu-eq88,1991A&A...242..286M}.  Furthermore, aerodynamic concentration is considered as a mechanism to increase densities locally for enhanced coagulation \\citep{Cuzzi_concentration,1997Icar..128..213K}.  However, most of these processes rely on a diversity in the aerodynamic properties of particles: Settling or radially drifting grains will only induce enhanced collision rates if not all grains settle or drift at the same rates.  Turbulent eddies will only enhance collision rates if different grains couple to different eddie scales. \\citet{1997LPI....28.1517W} reports that starting from 10$^{-4}$\\,cm grains, a hight central density of settled dust can still be achieved even if the fractal structure of the grains is taken into account.  Coagulation may proceed on quite different paths if the basic mechanism of coagulation is different from what is normally assumed. An interesting idea was put forward by \\citet{nu94} and \\citet{nu95} who studied aggregation of magnetic grains in laboratory experiments. They found that this accretion process proceeds rapidly forming large networks of linear aggregates.  \\citeauthor{nu94} suggested that such networks may be important for grain aggregation in the solar nebula. One way in which this mechanism could change the course of aggregation is to turn a cluster-cluster aggregation into a particle-cluster aggregation process.  If the networks of magnetic grains can form quickly enough, they may eventually provide a large fraction of the available collisional cross-section in the nebula.  The main aggregation process for non-magnetic grains would then be the collision of a grain with such a network.  This mechanism would be physically similar to a cluster-particle aggregation process which produces aggregates with fractal dimension $D_f=3$ \\citep{Ball84}, leading to more compact structures quickly.  In this article, we will present the foundations for a comprehensive investigation of magnetic dust aggregation from a theoretical and an experimental point of view.  We start by reviewing the evidence for magnetic grains in the solar nebula.  We will then focus on the physical implications of dipolar interaction between individual dust grains, i.e. enhanced collisional cross sections and accelerated aggregation dynamics.  We also present a new numerical tool we devised for the investigation of dust aggregation including a detailed treatment of grain-grain interaction through mechanical contacts, and dipolar (magnetic) forces.  We carried out numerical simulations of 2-body collisions with magnetic grains in order to test the code and compare the results with the analytically derived cross sections.  Two subsequent papers will communicate our experimental and numerical work of magnetic grain growth on the basis of the material presented here. ",
        "conclusions": "The material presented in this paper lays the foundation for a detailed analysis of the effects of magnetic aggregation for the accumulation of solid matter in the early solar system. Generalizing previous work by \\cite{hn99,hn00a}, we have derived an analytical expression for the collisional (capture) cross section of magnetic dust grains. As pairs of dipolar magnetic particles tend to line up in pole-to-pole configuration, aggregates of such particles ``preserve'' the magnetic moment of their constituents, thus displaying ``accretional remanence''. The combination of enhanced cross section and accretional remanence is responsible for accelerated growth dynamics.  This fact is quantitatively described by a dynamic exponent, $z$, such that the temporal dependence of the mean aggregate mass, $m$, takes on the form of a power law, i.e., $m\\propto t^z$.  We have shown that the dynamic exponent takes on a value $z \\approx 3.2$ for magnetic aggregation, whereas $z = 2$ for aggregation processes with non-magnetic particles.  Part of our theoretical approach has been validated numerically as a means for introducing a new coagulation code which we devised for the detailed analysis of magnetic coagulation. Its architecture allows the study of systems of particles interacting via long-range magnetostatic coupling as well as through short-range elastic forces. Aggregate restructuring is included. We provide evidence that such an approach -- although computationally rather costly - is required for physically realistic modelling.  A forthcoming paper will present extensive numerical studies of magnetic aggregation.   {\\bf Acknowledgements:} We would like to thank Tijmen van de Kamp for the development of Grainview, an interactive tool to visualize aggregate simulations; Vincent Icke and Kees Dullemond for discussons about how to implement grain rotation, Josef Sch\\\"ule for help with optimizing the code.  CD acknowledges financial support from Pioneer grant 600-78-333.  HN acknowledges financial support from the Studienstiftung des Deutschen Volkes and from DFG grant GL~142/13-1.  \\appendix"
    },
    "0201/astro-ph0201446_arXiv.txt": {
        "abstract": "A $\\gamma$-ray line production calculation in astrophysics depends on i) the composition and energy source spectrum of the energetic particles, ii) the propagation model, and iii) the nuclear cross sections.  The main difficulty for model predictions and data interpretation comes from the fact that the spectrum of the particles which actually interact in the ISM -- the \\textit{propagated spectrum}, is not the same as the \\textit{source spectrum} coming out of the acceleration site, due to energy-dependent energy losses and nuclear destruction.  We present a different approach to calculate $\\gamma$-ray line emission, based on the computation of the total number of photons produced by individual energetic nuclei injected in the interstellar medium at a given energy. These photon yields take into account all the propagation effects once and for all, and allow one to calculate quickly the $\\gamma$-ray line emission induced by energetic particles in any astrophysical situation by using directly their source spectrum.  Indeed, the same photon yields can be used for any source spectrum and composition, as well as any target composition.  In addition, these photon yields provide visual, intuitive tools for $\\gamma$-ray line phenomenology. ",
        "introduction": "High energy astrophysics is experiencing a considerable development, notably through the operation of a new generation of gamma-ray observatories, both on the ground and onboard satellites.  The analysis of gamma-ray emission from compact and diffuse sources is one of the most valuable ways to study high energy processes in the universe, and it is expected to put stronger and stronger constraints on the models in the near future.  Among the processes of interest, the production of gamma-ray lines by energetic particles (EPs) interacting in the interstellar medium (ISM) has received increased attention in the last few years, in connection with the data of the Compton Gamma-Ray Observatory as well as the forthcoming satellite INTEGRAL. In addition to the study of the high energy sources, the study of EP interactions is important for the understanding of the EPs themselves, of which the Galactic cosmic rays (GCRs) are one of the main components.  Information about the acceleration sites and processes is also provided by the determination of the EP energy spectrum and chemical composition, which can be derived, in principle, from the measurement of gamma-ray line ratios and profiles.  However, the information contained in the gamma-ray observational data relates to the spectrum and composition of the \\textit{propagated} particles (i.e. those who actually interact in the ISM), not the \\textit{source} particles.  The difference arises from the fact that the EPs experience various types of interactions while they `propagate' from their source to the place where they produce gamma-ray lines.  In particular, they lose energy through Coulombian interactions in a way which depends on both their energy and chemical nature, so that the propagated population of EPs is not identical to the source population, freshly coming out of the acceleration process.  Ideally, one would divide the whole process into three successive stages (e.g. Parizot and Lehoucq, 1999): particle acceleration, propagation and interaction, with the gamma-ray production arising during the last stage only.  In reality, of course, the reactions leading to gamma-ray production occur all the time, from the injection of the EPs into the ISM until they have slowed down to energies below the interaction thresholds.  It is therefore necessary to sum the contributions of all the instants following acceleration (the gamma-ray emission occurring during acceleration itself can usually be neglected, except for the rarest, highest energy particles, which spent a long time in the accelerator).  In steady state situations, this is equivalent to calculating the equilibrium distribution of EPs and integrating the relevant cross sections over this so-called propagated distribution.  From a technical point of view, the most difficult part consists in calculating the propagated spectrum, taking into account the energy losses and the energy-dependent escape time of the particles out of the confinement region, where the gamma-ray production is evaluated.  This is what prevents a straightforward calculation of the expected gamma-ray line fluxes from the knowledge of the nuclear cross sections and the EP source distribution.  Therefore, we propose here to work out this step once and for all, in the case of a steady state and a thick target, by calculating the integrated effect of energy losses on individual particles injected at any energy in the ISM. In Sect.~\\ref{sec.universality}, we shall justify the fact that the `propagation step' in a standard $\\gamma$-ray line emission calculation can be `factorized out' and calculated separately, independently of the EP source spectrum and composition.  This is true if the metallicity of the propagation medium does not exceed several tens of times the solar metallicity, as in most of the astrophysically sensible situations.  As a result, we shall obtain the absolute $\\gamma$-ray yields of energetic nuclei as a function of their initial energy, from which the $\\gamma$-ray line emission induced by EPs of any spectrum and composition can be straightforwardly calculated.  In addition to making $\\gamma$-ray line calculations much easier, these absolute yields (or particle efficiencies for $\\gamma$-ray line production) considerably help phenomenological interpretation of the observational data, as these yields only need to be convolved with the EP \\textit{source} spectra, rather than \\textit{propagated} ones. ",
        "conclusions": "In this paper, we have presented an easier way to calculate gamma-ray line emission from energetic particle interactions in the ISM. It is based on a simple mathematical transformation whose physical interpretation has been given and which allows one to work with the source spectrum of the EPs rather than the propagated spectrum.  Therefore, one does not need to worry about energy-dependent and nucleus-dependent energy losses of the particles, nor about their nuclear destruction in-flight, as they are taken into account once and for all through the calculation of absolute photon yields.  The latter are the number of photons produced in each of the nuclear de-excitation lines by a given nucleus injected in the ISM at a given initial energy.  These photon yields have been given here for various projectiles contributing to the main $^{12}$C, $^{14}$N, $^{16}$O, $^{20}$Ne and $^{56}$Fe de-excitation lines.  Numerical tables and electronic versions of the results are available from the authors upon request.  These photon yields are to be used instead of the nuclear excitation cross-sections, and might be thought of as `effective cross-sections' taking into account the specific effects of particle propagation in the ISM. They can be used to calculate the $\\gamma$-ray line emission induced by EPs with any spectrum and any composition in any medium with a metallicity lower than a few tens of the solar metallicity.  In addition to simplifying the calculation of gamma-ray line emission, the individual EP gamma-ray yields also provide a direct, intuitive tool to analyze gamma-ray line data from a phenomenological point of view, and construct an EP source spectrum and composition which could reproduce the intensity of the gamma-ray emission and the various line ratios.  The results presented here correspond to a thick target model, which is relevant to most astrophysical situations for EPs of energy lower than a few hundreds of MeV/n.  However, the same formalism can be used to calculate the total EP photon yields in a target with any escape length, be it energy-dependent or not.  Once these yields have been calculated once, they can be used in any situation with the same escape length, for any particle spectrum and any EP and target compositions."
    },
    "0201/astro-ph0201387_arXiv.txt": {
        "abstract": "Theoretical predictions of Red Giant Branch stars' effective temperatures, colors, luminosities and surface chemical abundances are a necessary tool for the astrophysical interpretation of the visible--near infrared integrated light from unresolved stellar populations, the Color-Magnitude-Diagrams of resolved stellar clusters and galaxies, and spectroscopic determinations of red giant chemical abundances. On the other hand, the comparison with empirical constraints provides a stringent test for the accuracy of present generations of red giant models.  We review the current status of red giant stars' modelling, discussing in detail the still existing uncertainties affecting the model input physics (e.g., electron conduction opacity, treatment of the superadiabatic convection), and the adequacy of the physical assumptions built into the model computations.  We compare theory with several observational features of the Red Giant Branch in galactic globular clusters, such as the luminosity function \\lq{bump}\\rq, the luminosity of the Red Giant Branch tip and the envelope chemical abundance patterns, to show the level of agreement between current stellar models and empirical data concerning the stellar luminosities, star counts, and surface chemical abundances. ",
        "introduction": "The Red Giant Branch (RGB) is one of the most prominent and well populated features in the Color-Magnitude-Diagram (CMD) of stellar populations with ages larger than about $1.5 - 2$ Gyr. The theoretical modelling of RGB stars plays therefore a wide ranging role, involving various fields of galactic and extragalactic astrophysics.  Since RGB stars are cool, reach high luminosities during their evolution, and their evolutionary timescales are relatively long, they provide a major contribution to the integrated bolometric magnitude and to integrated colors and spectra at wavelengths larger than about 900 -- 1000 nm of old distant, unresolved stellar populations (e.g. Renzini \\& Fusi-Pecci~1988; Worthey~1994). A correct theoretical prediction of the RGB spectral properties and colors is thus of paramount importance for interpreting observations of distant stellar clusters and galaxies using population synthesis methods, but also for determining the ages of resolved globular and open clusters, by means of isochrone fitting techniques.  The $I$-band brightness of the tip of the RGB (TRGB) provides a robust standard candle, very much independent of the stellar age and initial chemical composition, which can allow to obtain reliable distances out to about 10 Mpc using $HST$ observations (e.g., Lee, Freedman \\& Madore~1993 -- LFM93). Due to the lingering uncertainties on the empirical determination of the TRGB brightness zero point, RGB models provide an independent path to the calibration of this important standard candle (Salaris \\& Cassisi~1997). Also the theoretical calibration of the luminosity of Horizontal Branch (HB) stars and their RR Lyrae population (whose parallax-based distances still show a very large error bar; see, e.g., Groenewegen \\& Salaris~1999) is dependent on the correct modelling of the previous RGB phase, since HB luminosities (like the TRGB ones) are determined by the value of the electron degenerate He-core mass ($M_{core}^{He}$) at the end of the RGB evolution.  Predicted evolutionary timescales along the RGB phase play a fundamental role in the determination of the initial He abundance of globular cluster stars through the R parameter (number ratio between HB stars and RGB stars brighter than the HB at the RR Lyrae instability strip level; see, e.g. Iben~1968a, Sandquist~2000, Zoccali et al.~2000), while an accurate modelling of the mixing mechanisms efficient in the RGB stars is necessary to correctly interpret spectroscopic observations of their surface chemical abundance patterns.  These various applications of RGB stellar models to fundamental astrophysical problems crucially rely on the ability of theory to predict correctly:  \\noindent -- the CMD location (in $T_{\\rm eff}$ and color) and extension (in brightness) of the RGB as a function of the initial chemical composition and age;  \\noindent -- the evolutionary timescales (hence the relative numbers of stars at different luminosities) all along the RGB;  \\noindent -- the RGB stars' physical and chemical structure, as well as their evolution with time.  The main goal of this review is to discuss the existing uncertainties in theoretical RGB models, and to assess the reliability of their predictions.  In \\S 2 we present an outline of RGB stellar evolution, while in \\S 3 and \\S 4 fundamental properties of the CMD and luminosity functions of RGB models are reviewed. \\S 5 analyzes the input physics currently used for computing RGB models, and how related uncertainties affect the outcome of the calculations; \\S 6, \\S 7 and \\S 8 discuss observational tests for the accuracy and adequacy of RGB models by employing CMDs, spectroscopic observations, and luminosity functions. \\S 9 reviews the use of the TRGB as standard candle, and conclusions follow in \\S 10. ",
        "conclusions": "RGB theoretical models play a key role in the interpretation of various astrophysical observations. We have discussed in detail the theoretical uncertainties still affecting the models, paying particular attention to the predictions of colors, luminosities, evolutionary timescales and surface chemical abundances. A comparison of various theoretical RGB models with a large body of diverse empirical data has allowed us to discuss not only the accuracy of the adopted input physics, but also the adequacy of the assumptions built into canonical RGB models.  Large uncertainties still exist in the predictions of $T_{\\rm eff}$ and colors, due mainly to the treatment of superadiabatic convection, boundary conditions and color transformations employed; they can however be minimized by a suitable calibration of the models on empirical data. The prediction of the absolute value of the He-core mass at the He-flash -- which determines the TRGB and HB brightness -- suffers from residual uncertainties mainly related to the determination of accurate electron conduction opacities in the relevant physical regime.  Extra mixing processes not included in canonical RGB models seem to be required to help explaining some of the chemical abundance patterns observed at the photosphere of RGB stars; however, independent empirical constraints, and arguments from stellar nucleosynthesis require that these processes should not be able to affect appreciably the evolutionary timescales, He-core masses and CMD location obtained from canonical models."
    },
    "0201/astro-ph0201452_arXiv.txt": {
        "abstract": "We review the progress of research on intracluster planetary nebulae (IPN).  In the past five years, hundreds of IPN candidates have been detected in the Virgo and Fornax galaxy clusters and searches are also underway in poorer galaxy groups.  From the observations to date, and applying the known properties of extragalactic planetary nebulae, the intracluster light in Virgo and Fornax: 1) is significant, at least 20\\% of the total cluster stellar luminosity, 2) is elongated in Virgo along our line of sight, and 3) may derive from lower-luminosity galaxies, consistent with some models of intracluster star production.  A fraction of IPN candidates are not true IPN, but emission-line sources of very large observed equivalent width ($\\geq 200$ \\AA). The most likely source for these contaminating objects are Lyman-$\\alpha$ galaxies at z $\\approx$ 3.1. Follow-up spectroscopy of the IPN candidates will be crucial to discriminate against high redshift galaxies and to derive the velocity field of the intracluster stellar population. ",
        "introduction": "Intracluster stars, stars between the galaxies in a galaxy cluster, are an excellent probe of galaxy and galaxy cluster evolution. They preserve a record of a galaxy cluster's dynamical history, and are a sensitive measure of the poorly understood processes of galactic mergers, cluster accretion, and tidal-stripping that occur in galaxy clusters (cf.~Dressler 1984).  The original concept of intracluster stars was first proposed over fifty years ago (Zwicky 1951), but actually observing intracluster starlight has been extremely difficult. Typically, the surface brightness of intracluster starlight is less than 1\\% that of the brightness of the sky background, and measurements of this luminosity must contend with the problems presented by scattered light from bright objects and the contribution of discrete sources below the detection limit. These difficulties have made research into intracluster starlight slow and uncertain, though solid detections have been made (for reviews, see V{\\'i}lchez-G{\\'o}mez 1999; Feldmeier 2000).  An alternative way to study intracluster light is to search for it in nearby galaxy clusters.  Here, it is possible to search for luminous individual intracluster stars, and gain more detailed information on the distribution, metallicity and velocities of intracluster stars than is possible from surface brightness measurements.  Planetary nebulae (PN) are one such tracer of the intracluster starlight.  The history of intracluster planetary nebulae (IPN) research begins a decade ago with the first PN survey of the Virgo cluster (Jacoby, Ciardullo, \\& Ford 1990; hereafter JCF).  In this survey of elliptical galaxies, JCF found 11 PN that were much brighter than the expected [O~III] $\\lambda$ 5007 planetary nebulae luminosity function (PNLF) cut-off magnitude.  JCF attempted to explain these ``overluminous'' PN with a number of hypotheses, but none was entirely satisfactory.  The next step involved spectroscopic follow-up of objects from the JCF survey.  During a radial velocity survey of PN in the Virgo elliptical galaxy M~86, Arnaboldi \\etal (1996) found that 16 of the 19 detected PN velocities were consistent with the galaxy's mean velocity (v$_{radial}$ = -227 \\kms).  The other three planetaries had mean velocities of $\\sim 1600$ \\kms, more consistent with the Virgo cluster's mean velocity.  Arnaboldi \\etal (1996) argued convincingly that these objects were intracluster planetary nebulae, and it is here that the term first enters the literature. Almost simultaneously, the first search for IPN candidates in the Fornax cluster was published (Theuns \\& Warren 1997), and more detections of IPN candidates in Virgo quickly followed (M\\'endez \\etal 1997; Ciardullo \\etal 1998; Feldmeier, Ciardullo, \\& Jacoby 1998).  Additionally, Ferguson, Tanvir and von Hippel (1998) detected intracluster red giant stars (IRGs) in the Virgo cluster using {\\sl HST}.  This provided independent confirmation of the IPN results, and allowed for a potential comparison between the two detection methods.  With the help of large CCD mosaics, samples of hundreds of IPN have now been gathered in the Virgo and Fornax clusters (Fig 1., Feldmeier 2000; Arnaboldi \\etal 2002).  IPN candidates are easily identified as stellar sources that appear in a deep [O~III] $\\lambda$ 5007 image, but completely disappear in an image through a filter that does not contain the [O~III] line.  Originally, searching for such objects was done manually, but now is done through automated methods (Feldmeier 2000; Arnaboldi \\etal 2002) \\begin{figure} {\\vskip100pt} \\caption{Images of portions of the Virgo (top) and Fornax (bottom) galaxy clusters, with the locations of IPN and other intracluster star detections marked. There are 319 IPN candidates found in Virgo Fields 1-8, and 138 candidates in the Fornax fields.  Virgo's subclumps A \\& B are centered at the the top and bottom of the Virgo image.} \\end{figure} ",
        "conclusions": ""
    },
    "0201/astro-ph0201178_arXiv.txt": {
        "abstract": "In order to investigate the properties of the prominent tidal debris feature extending to the northeast of the compact group of galaxies Seyfert's Sextet, we analyzed multi-band ($U$, $B$, $V$, $VR$, $R$, $I$, $J$, $H$ and $K^{\\prime}$) photometric imaging data and obtained the following results: 1) The radial surface brightness distribution of this tidal debris in Seyfert's Sextet (TDSS) in each band appears to be well approximated by an exponential profile. 2) The observed $B-V$ color of TDSS is similar to those of dwarf elliptical galaxies in nearby clusters. 3) Comparing the spectral energy distribution (SED) of TDSS with theoretical photometric evolution models and with the SED of the stars in the outer part of HCG 79b, we find that its SED is comparable to that of a $\\sim$ 10 Gyr-old stellar population with solar metallicity, similar to the stellar population in the outer part of HCG 79b. This suggests that TDSS consists of stars that may have been liberated from HCG 79b by strong galaxy interactions, not a pre-existing dwarf galaxy previously thought. ",
        "introduction": "Although dwarf galaxies are the most numerous extragalactic objects in the nearby Universe (e.g., Ferguson \\& Binggeli 1994; Binggeli, Sandage \\& Tammann 1988; Mateo 1998), it seems unclear how they are related in origin to typical large (i.e., $L^{*}$) galaxies. Dwarf galaxies could be formed through the same formation mechanism as that of large galaxies; e.g., gravitational collapse of protogalactic gas clouds (Dekel \\& Silk 1986; White \\& Frenk 1991; Frenk et al. 1996; Kauffmann, Nusser \\& Steinmetz 1997). However, it is known that dwarf elliptical galaxies (dEs) apparently belong to a different class from normal large ellipticals (Es) in the fundamental plane (e.g., Kormendy 1985), suggesting that the formation and/or evolution processes of dwarfs may not always be the same as those of larger ellipticals.  It has been argued from an observational view point that dwarf galaxies may be formed by galaxy collisions because there appears to be morphological evidence for dwarf galaxies in the tidal tails of interacting galaxies (Zwicky 1956; Schweizer 1978; Duc et al. 2000); i.e., gas-rich dwarf irregular galaxies, can be made out of stellar and gaseous material pulled out into intergalactic space by tidal forces from the disks of colliding parent galaxies.  This possibility has been recently reinforced by  a number of pieces of observational evidence (Schweizer 1982;  Bergvall \\& Johansson 1985; Schombert, Wallin \\& Struck-Marcell 1990;  Mirabel, Lutz \\& Maza 1991; Mirabel, Dottori \\&  Lutz 1992; Duc \\& Mirabel 1994, 1998; Duc et al. 2000; Yoshida, Taniguchi \\& Murayama 1994; Braine et al. 2000; Weilbacher et al. 2000).  Also, Hunsberger et al. (1996; 1998) find an excess of dwarf galaxies in compact groups of galaxies apparently caused by interactions among group members. Such formation of tidal dwarf galaxies (TDGs) has also been demonstrated by numerical simulations of merging/interacting galaxies (Barnes \\& Hernquist 1992; Elmegreen, Kaufman \\& Thomasson 1993). Therefore, tidal formation seems to potentially be an important formation mechanism for dwarf galaxies (Okazaki \\& Taniguchi 2000 and references therein).  One famous tidal debris system extends to the northeast of Seyfert's Sextet (hereafter SS). SS is one of the most famous, as well as densest, compact groups of galaxies (Seyfert 1948a, 1948b; see for a review, Raba\\c{c}a 1996). This group is also a Hickson compact group (hereafter HCG) of galaxies, HCG 79 (Hickson 1982; 1993). Many subsequent studies of SS have mentioned that the galaxies in SS appear to show morphologically and dynamically peculiar properties  (Sulentic \\& Lorre 1983; Rubin, Hunter, \\& Ford 1991;  Bettoni \\& Fasano 1993; Mendes de Oliveira \\& Hickson 1994;  Bonfanti et al. 1999; Nishiura et al. 2000a). Hickson himself regarded SS as a galaxy quartet (HCG 79a, 79b, 79c, and 79d).  A fifth component, HCG 79e, was found to be a redshift-discordant galaxy that is believed to  have no physical relation to SS (Hickson 1992).   A sixth object, or more precisely the north-eastern optical fuzz, is now considered likely to be tidal debris associated with the morphologically peculiar galaxy HCG 79b (Rubin et al. 1991; Williams et al. 1991; Mendes de Oliveira \\& Hickson 1994; V\\'{\\i}lchez \\& Iglesias-P\\'{a}ramo 1998).  Current X-ray observations are not sensitive enough to detect any X-ray emission that may originate from intragroup gas that might be present in SS (Pildis, Bregman, \\& Evrard 1995; Ponman et al. 1996).  However, Sulentic \\&  Lorre (1983) and Nishiura et al. (2000b) detected a faint optical envelope around SS that is plausibly composed of stars tidally liberated from the galaxies in SS, plus Williams, McMahon \\& van Gorkom (1991) found extended  H{\\sc i} emission.  These observations suggest that SS is a physically real compact group.  In this paper, we present results of our photometric study of TDSS. Since this tidal debris system in SS (hereafter TDSS) is morphologically similar to other tidal debris, such as Arp 105S and Arp 245N (Braine et al. 2000), we will also compare the photometric properties of TDSS with these two tidal debris. Throughout this paper we adopt a distance to SS of 44 Mpc determined using the mean  recession velocity of HCG 79a, 79b, 79c, and 79d referenced to the  galactic standard of rest, $V_{\\rm GSR}$ = 4449 km s$^{-1}$ (de Vaucouleurs et al. 1991), and a Hubble constant, $H_{0}$ = 100 km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "From multi-band photometry of the prominent tidal debris feature to the northeast of Seyfert's Sextet we have obtained the following results: \\begin{enumerate} \\item The surface brightness profile of this tidal debris in Seyfert's Sextet (TDSS) in each band shows an approximately exponential profile. \\item The observed $B-V$ color of TDSS is redder than those of tidal dwarf galaxies, e.g., Arp 105S and Arp 245N (Braine et al. 2000). \\item Comparing the spectral energy distribution (SED) of TDSS with theoretical photometric evolution models, we find that its SED is comparable to that of a stellar population with age $\\sim$ 10 Gyr and with higher metallicity than the average values found in tidal dwarf galaxies. \\item Comparing SEDs, we find that the SED of TDSS is similar to that of the stars in the outer part of HCG 79b. \\item TDSS seems to consist primarily of stars liberated from HCG 79b by strong galaxy interaction. \\end{enumerate}  We conclude that TDSS is simply a passive tidal feature like many others in interacting galaxy systems. We also conclude that there is no indication that TDSS will evolve into a star-forming tidal dwarf galaxy similar to those previously studied. This, however, indicates that another type of forming dwarf galaxy without secondary star formation through the galaxy interaction.  \\vspace{1ex}  We would like to thank the staff members of the Okayama Astrophysical Observatory, the KISO observatory and the UH 2.2 m telescope for their kind assistance during our observations. We thank J. W. Sulentic for his HST observation of Seyfert's Sextet, and an anonymous referee for several useful comments which helped improve the paper. We also thank Ichi Tanaka for his kind help during our observations, Richard Wainscoat and Shinki Oyabu  for their useful comments on photometric calibration,  and Daisuke Kawata for useful discussions.  A part of this work was done when YT was a visiting astronomer  at the IfA, University of Hawaii. YT would like to thank Rolf-Peter  Kudritzki and Bob McLaren for their warm hospitality. YS thanks the Japan Society for Promotion of Science (JSPS) Research Fellowships for Young Scientist. This work was supported in part by the Ministry of Education, Science, Sports and Culture in Japan under Grant Nos. 07055044, 10044052, and 10304013."
    },
    "0201/astro-ph0201514_arXiv.txt": {
        "abstract": "The temperature of the diffuse, photo-heated intergalactic medium (IGM) depends on its reionization history because the thermal time scales are long. The widths of the hydrogen \\lya absorption lines seen in the spectra of distant quasars that arise in the IGM can be used to determine its temperature. We use a wavelet analysis of the \\lya forest region of quasar spectra to demonstrate that there is a relatively sudden increase in the line widths between redshifts $z\\approx 3.5$ and 3.0, which we associate with entropy injection resulting from the reionization of \\Hep. The subsequent fall-off in temperature after $z\\approx 3.5$, is consistent with a thermal evolution dominated by adiabatic expansion. If, as expected, the temperature also drops rapidly after hydrogen reionization, then the high temperatures inferred from the line widths before \\Hep\\ reionization imply that hydrogen reionization occurred below redshift $z=9$. ",
        "introduction": "Neutral hydrogen in the intergalactic medium (IGM) along the line of sight to quasars at redshifts $z\\le 6$ produces hundreds of \\lya absorption lines.  The fact that not all flux is absorbed (i.e., the absence of a \\lq Gunn-Peterson\\rq~ trough, Gunn \\& Peterson 1965) requires that the universe be ionized to a level far higher than can be attributed to residual ionization from recombination. At lower redshifts, $z\\la 3$, {\\em observed} stars and quasars produce enough ionizing photons to explain the high levels of ionization, but the nature of the sources responsible for converting most of the IGM from neutral to ionized remain uncertain, as does the epoch of reionization (e.g., Barkana \\& Loeb 2001).  The observed mean flux decrement $D_A$ blueward of the quasar's \\lya emission line increases with the redshift of the quasar, both because the intensity of the ionizing background radiation decreases above $z=4$ (e.g., McDonald \\& Miralda-Escud\\'e 2001) and because the mean density $\\bar\\rho$ of the universe -- and hence the neutral fraction $x\\propto \\bar\\rho T^{-0.7}/\\Gamma_\\h$ for fixed values of the photo-ionization rate $\\Gamma_\\h$ and temperature $T$ -- increases. Recently, Becker et al.\\ (2001) and Djorgovski et al.\\ (2001) observed a sudden increase in $D_A$ in the spectra of redshift $z\\sim 6$ quasars discovered by the {\\sc SLOAN} digital sky survey. Such a sharp rise has been predicted to mark the transition associated with a sudden epoch of reionization (e.g., Cen \\& Ostriker 1993; Gnedin 2000). Similarly, a sudden increase in the \\Hep\\ opacity has been detected around $z\\sim 3$ (Reimers et al.\\ 1997; Heap et al.\\ 2000; Kriss et al.\\ 2001), associated with helium reionization.  Another way to study the reionization history of the IGM is to investigate its thermal evolution. Because its cooling time is long, the low-density IGM retains some memory of when and how it was reionized (e.g., Miralda-Escud\\'e \\& Rees 1994; Hui \\& Gnedin 1997; Haehnelt \\& Steinmetz 1998).  The combined effects of photo-ionization heating and adiabatic expansion introduce a tight temperature-density relation in the unshocked IGM, which can be approximated by a power-law $T=T_0 (\\rho/\\bar\\rho)^{\\gamma-1}$ for densities around the cosmic mean (Hui \\& Gnedin 1997). A change in these parameters influences the shapes of the \\lya lines, because thermal broadening and Jeans smoothing determine the line widths (e.g., Theuns, Schaye \\& Haehnelt 2000). Schaye et al.\\ (1999) used hydrodynamical simulations to demonstrate that one can accurately calibrate the relation between the minimum line width ($b$) as a function of column density ($N_\\h$) on the one hand, and the underlying $T-\\rho$ relation on the other. Schaye et al.\\ (2000) applied the method to observations in the redshift range 2.0-4.5 and found that $T_0$ peaks at $z\\sim 3$, which they interpreted as evidence for the reionization of \\Hep. Ricotti, Gnedin \\& Shull (2000) used pseudo hydrodynamical simulations and found a similar temperature increase, albeit only at the 0.5$\\sigma$ level. McDonald et al.~(2001) found no evidence for temperature evolution, but their analysis neglected the important temperature dependence of Jeans smoothing. Finally, the analysis of Zaldarriaga et al.\\ (2001) neglected hydrodynamical effects all together.  Here, we provide new evidence for a relatively sudden increase in $T_0$ between redshifts $z\\approx 3.5$ and 3.0, using a new method based on a wavelet decomposition of the absorption spectrum. We then use the measured values of $T_0$ at higher redshift to constrain the epoch of hydrogen reionization $z_{\\rm H}$ and find that the data require $z_{\\rm H} \\le 9$ for any reasonable value ($T_0\\la 6\\times 10^4~\\K$) of the hydrogen reionization temperature.  \\newpage ",
        "conclusions": ""
    },
    "0201/astro-ph0201272_arXiv.txt": {
        "abstract": "An optical photometric observation with the Beijing-Arizona-Taiwan-Connecticut (BATC) multicolor system is carried out for the central region of the nearby cluster of galaxies, Abell 2634. From the $2k \\times 2k$ CCD images with fourteen filters which covers a range of wavelength from 3600 \\AA ~to 10000 \\AA,  5572 sources are detected down to $V \\sim 20$ mag in a field of 56' $\\times$ 56' centered on this regular cluster of galaxies. As a result, we achieved the spectral energy distributions (SEDs) of all sources detected. There are 178 previously known galaxies included in our observations, 147 of which have known radial velocities in literature. After excluding the foreground and background galaxies, a sample of 124 known members is formed for an investigation of the SED properties. The comparison of observed SEDs of the early-type member galaxies with the template SEDs demonstrates the accuracy and reliability of our photometric measurements.  Based on the knowledge of SED properties of member galaxies, we performed the selection of faint galaxies belonging to Abell 2634. It is well shown that the color-color diagrams are powerful in the star/galaxy separation. As a result, 359 faint galaxies are selected by their color features.  The technique of photometric redshift and color-magnitude correlation for the early-type galaxies are applied for these faint galaxies, and a list of 74 faint member galaxies is achieved. On the basis of the  newly-generated  sample of member galaxies, the spatial distribution and color-magnitude relation of the galaxies in core region of Abell 2634 are discussed. There exists a tendency that the color index dispersion of the early-type members is larger for the outer region, which might reflect some clues about the environmental effect on the evolution of galaxies in a cluster. ",
        "introduction": "It is widely appreciated that the observations of galaxy clusters can provide useful constraints on the theories of formation and evolution of large-scale structure, since they are the largest gravitationally bound systems in the Universe. A great volume of observational data has been achieved for the galaxies belonging to a cluster, particularly for the nearby clusters of galaxies. The rich cluster of galaxies, Abell 2634, is one example. Extensive observational efforts are being made to expand the kinematic database for Abell 2634 because it  has many intriguing observational features. First, its central galaxy,  NGC 7720, classified as cD first by Matthews, Morgan, \\& Schmidt (1964) and then classified as D galaxy by Dressler (1980a), was found to have a  companion galaxy, namely NGC 7720A, with a projected distance of 7.2  $h_{75}^{-1}$ kpc and a velocity difference of about 1000 km s$^{-1}$.  Second, a surprisingly large negative peculiar velocity for Abell 2634 was derived when the fundamental plane (FP) relation   $D_n - \\sigma$ (Dressler et al. 1987) was sketchily extended to estimate the distance of Abell 2634  (Lucey et al. 1991), which is challengeable for the universality of the FP and Tully-Fisher (TF) relations. With the effort for several years by Lucey et al. (1997) and Scodeggio, Giovanelli \\& Haynes (1997), larger samples of member galaxies with more accurate photometric measurements (including $I$ band observations) are constructed. Abell 2634 was found to have a negligibly small peculiar velocity with respect to the cosmic microwave background reference frame. The FP  distance estimate of Abell 2634 is in good agreement with that given by the TF relation.  Additionally, the prototypical wide-angle tailed (WAT) radio source, 3C465, is found to coincide with the central cD galaxy, NGC 7720 (Griffen 1963; Eilek et al. 1984; O'Donoghue, Owen \\& Eilek 1990). A study of this cluster by Pinkney et al.(1993) suggested that the large scale turbulent gas motions fueled by a cluster-subcluster merging process might be responsible for the observed bending of radio tails.  To investigate issues related to the kinematics of NGC 7720, Scodeggio et al. (1995; hereafter SSGH) carried out the multi-fiber spectroscopy and 21 cm line observations, and particularly analyzed the structure, kinematics, and morphological segregation of galaxies within the central half degree region. There are 99 member galaxies bright than $m_v \\sim 16.0$ in the inner $0.^{\\circ}5$ region, which is called the HDR (half degree region) sample. The completeness of the inner region of one square degree reached $m_{pg} \\sim 16.0$. As a result, the morphological dependence of spatial and kinematic properties was found: the early-type galaxies appear to be a relaxed system, while the spirals have much larger velocity dispersion.  Studies of the luminosity function of cluster galaxies show that most intrinsically bright galaxies in clusters, $M_B < -16$, are giant ellipticals. Down to $m_v \\sim 18.5$ in Abell 2634, the majority of apparently faint member galaxies are still intrinsically bright. There should be many giant ellipticals within the apparent magnitude range of $16.0 < m_v < 18.5$. In order to enrich our understanding of the structure and dynamics of Abell 2634, these faint member galaxies should be taken into account. However, the spectroscopic observations are not available yet for a great number of galaxies fainter than 16.0 mag in Abell 2634. The multicolor optical photometric observation therefore becomes a common alternative means to study some properties of the member galaxies, such as the color-magnitude relation (Bower, Lucey \\& Ellis 1992), Butcher-Oemler effect (Butcher \\& Oemler 1978) and morphology-density relation (Dressler 1980b).  Due to a large field of view, the Schmidt Telescope (1-meter class) telescope equipped with a large format CCD, is one of most suitable facilities for obtaining photometric data for the nearby clusters of galaxies. The multicolor photometry can provide the spectral energy distributions (SEDs) of all the objects within the field of view. Comparing the integration times required for spectroscopy to a certain depth on a large telescope, the amount of observing times for multicolor photometry are much shorter and therefore more readily scheduled. The Beijing-Arizona-Connecticut (BATC) multicolor photometric survey is designed to obtain the SED information of cluster galaxies and other types of objects (Fan et al. 1996). This paper will present the multicolor optical  photometry of the central region of Abell 2634, using the 60/90 cm Schmidt Telescope of Beijing Astronomical Observatory (BAO) with 14 BATC filters covering a wide wavelength range from 3800 \\AA ~to nearly 10000 \\AA.  The comparison of observed SEDs for the previously known member galaxies with the template SEDs demonstrates the accuracy and reliability of our photometric measurements. Our multicolor photometry contains a large quantity of bright member galaxies with sufficiently small photometric errors, which allows a verification of color-magnitude correlation as well as a testing application of the technique of photometric redshift. After performing a reliable star/galaxy separation by the color-color diagrams, we carried out a membership selection based on the SED features of known member galaxies. The technique of photometric redshift and color-magnitude correlation have provided useful constraints on the membership selection. As a result, we cataloged the SEDs of the newly-selected faint members, based on which some properties of the enlarged sample of member galaxies can be addressed. We believe that current study will not only supplement the data bases of member galaxies in this well-studied cluster, but also extend the investigations on the spatial distribution and color-magnitude relation of the galaxies in the core of Abell 2634 to an unprecedented depth.  This paper is structured as follows. In Section 2, we will present our observations and data reduction, as well as the photometric catalogs of the objects (including the known member galaxies) in the central field of Abell 2634.  The analysis of the SEDs for the known member galaxies is shown in Section 3.  A SED selection of faint member galaxies is given in Section 4, and the spatial distribution and color properties of the enlarged sample of member galaxies is presented in Section 5. Finally, we will give a summary in Section 6.  The cosmological parameters $H_0=50$ km s$^{-1}$ Mpc$^{-1}$ and $q_0=0.5$ have been assumed.   \\section {Observations and data reduction} \\subsection{BATC observations}  Abell 2634 is a nearby ($z \\sim 0.03$) regular cluster of galaxies, classified by Abell (1958) as of richness class 1. Dressler (1980a) listed 132 galaxies in his catalog, and the fraction of early-type galaxies (E/S0) is about 63\\%. In the HDR sample of SSGH, the fraction of early-type galaxies reaches $\\sim$ 71\\%. This cluster appears to be at a distance of $\\sim$ 9000 km s$^{-1}$, projected on the main ridge of the Pisces-Perseus Supercluster (PPS) (see Fig.1 in Wegner, Haynes, \\& Giovanelli 1993). Its companion Abell 2666 is closely located at approximately 3 degree to the east, with a heliocentric redshift of $\\sim$ 8154 km s$^{-1}$ (Struble \\& Rood 1999).  Our multicolor optical photometry concentrates on a region of 56 $\\times$ 56 arcmin$^2$ centered on NGC 7720. The 69/90 cm f/3 Schmidt Telescope of BAO, located at Xinglong site with an altitude of 900m, was used to obtain the photometric observations. A Ford 2048 $\\times$ 2048 CCD camera was equipped at the prime focus of the telescope. The field of view was 0.95 square degree and the spatial scale was $1.67''$ per pixel. The details of the Schmidt Telescope, CCD camera and data-acquisition system can be found elsewhere (Fan et al. 1996; Yan et al. 2000). We make use of the BATC filter system which includes 16 intermediate-band filters covering the whole optical wavelength range from $\\sim$ 3000 \\AA ~to 10000 \\AA. The transmissions of these filters can been seen in Fig. 1 of Zhou et al. (2001).  Only 14 filters, from $b$ to $p$, were used for the observations of Abell 2634. The filter number, filter name, effective wavelength centroid and FWHM for each filter are listed in Table 1. In total, we have made more than 40 hours exposures, and obtained 214 images on the central region of Abell 2634. After a check of the image quality, 173 images with nearly 36 hours of exposure were selected to be combined.  The combined images cover an area defined by a right ascension range from $23^h 36^m 16^s$ to $23^h 40^m 23^s$, and a declination range from $26^{\\circ} 33' 30''$ to $27^{\\circ} 30' 01''$ (for equinox of 2000.0).  The seeing of the combined images were quite different from band to band, but the typical seeing of combined image for each filter band was about 5 arcsec.  Almost all nights were thought to be photometric by the observers. The standard  stars were observed  between  air-masses 1  to 2 for each program filter band. Normally, four to six filters were selected to make flux calibrations during each photometric night. The standard stars were put near the CCD center and only 300 $\\times$ 300 sub-images were taken for the standard stars to save readout time and disk space. The extinction coefficient and magnitude zero point obtained from standard stars were used for making the flux calibration on the BATC field images. The details of the observation for calibration are described in Yan et al. (2000). Table 1 also gives the parameters for our observations.   \\subsection{Data reduction}   The bias subtractions and dome flat-field corrections were done on the CCD images. The cosmic ray and bad pixel effect were corrected by comparing the images. The images were re-centered  and  position calibration was performed by using the Guide Star Catalog (GSC). The fluxes of intermediate-band filter images  were calibrated by Oke-Gunn standard stars (Oke \\& Gunn 1983; Fukugita et al. 1996).  As in other works on this image field, we are preparing the spectral energy distribution (SED) catalog of sources. The magnitudes were measured by the aperture photometry with detection threshold of $4\\sigma$ of the sky fluctuation per pixel. A minimum of 1 pixel above the threshold was required for detection. This might introduce a false detection, particularly for the nearby very bright stars. Since most of galaxies in the images are obviously  extended, we finally selected a fixed aperture of 5 pixels to do the photometry, which is large enough to make different seeing effects negligible. Although the resulting magnitudes given by photometry with fixed aperture are not the same as the total magnitudes of galaxies in literatures, this is a proper way to obtain the reliable color indices (i.e., relative SEDs) of all the sources. The photometry  program developed  by Bertin \\& Arnouts (1996), SExtractor (Ver2.1.6) namely, was used for the photometry. During extraction of the sources, a filter of $3\\times3$ ``all-ground'' convolution mask with FWHM = 2 pixels was used. Since the adopted aperture size is rather large, contamination of the light by nearby objects may sometimes not be negligible. The number of sources detected in the combined images are shown in Table 1.  We used four standard stars, namely, BD+17d4708, BD+26d2606, HD19445, and HD84937, in Oke \\& Gunn (1983) for the BATC flux calibration. The absolute fluxes of these stars were taken from Fukugita et al. (1996). The definitions of BATC magnitude are in the AB$_{\\nu}$ system of Oke \\& Gunn (1983):$ m_{\\rm BATC} = -2.5\\cdot {\\rm log} \\widetilde{F_{\\nu}} - 48.6. $ Once the aperture photometry on the images of standard stars was done, the extinction coefficients, $K$, and magnitude zero points, $C$, can be obtained by fitting the extinction curve of magnitude {\\it versus} airmass: $ m_{inst}-m_{BATC} = K\\cdot X + C, $ where $X$ is the airmass of the image.  $K$ and $C$ were derived by a median fitting of the data points with straight line. The procedures of flux calibration  are detailed in Zhou et al. (2001).  For $c$ and $o$ filters, we do not have the calibration images. Fortunately, we developed a method to calibrate the SEDs of objects for our large field multicolor photometric system, based on the the SED library (Zhou et al. 1999), which is the so-called model calibration. This method heavily improved the quality of the observations taken under not so perfect photometric condition. The SED calibration was done for our combined images. Then, we can derive the calibrated magnitudes in $c$ and $o$ bands from the information of the calibrated SEDs. For other filter bands, the magnitudes derived by ordinary standard calibrations were adopted for further analyses.  \\subsection{Photometric catalog}  We achieved the photometric magnitudes within 14 filter bands for 5572 sources in the central region (56' $\\times$ 56') of Abell 2634. This SED catalog will be electronically provided upon request. To cross-identify all the known galaxies within the our field, we made use of the NASA/IPAC Extragalactic Database (NED). As a result, 178 known galaxies were found to have counterparts in our SED catalog. The identification is unambiguous for a majority of these galaxies, according to their positional offsets and apparent magnitudes. The SEDs of these 178 galaxies are cataloged in Table 2, which structured as follows:  \\begin{tabular}{ll} Column 1: & Number of the galaxies, sorted by the R.A.in 2000 epoch.\\\\ Column 2,3: & R.A. and declination in 2000 epoch, in ``hhmmss.ss'' and ``ddmmss.ss'' \\\\ &mode, given by our photometric measurements. \\\\ Column 4: & Redshift, in km s$^{-1}$, provided by the NED.\\\\ Column 5: & Membership code: member galaxies are identified as ``m''; foreground \\\\ & galaxies are marked as ``f''; background galaxies are marked by ``b''; \\\\ & and ``?'' means possible members. See text for details. \\\\ Column 6-19: & Photometric magnitudes in $b - p$ filter bands. The value of 0.0 means\\\\ & no detection in this filter band. \\\\ \\end{tabular}  It should be noted that the photometric magnitudes given in Table 2 might be somewhat different from the total magnitudes of galaxies given in some catalogs, since we used a fixed aperture to do the photometry. What we attempted to obtain was the relative SEDs of objects in the central region of  Abell 2634.   We checked two kinds of errors respectively given by the photometric program SExtractor and by a comparison with the stellar spectral templates of Gunn \\& Stryker (1983). The statistic errors given by the SExtractor are smaller than the real measurement errors. We compared the errors using different images with the same filter. By separating stars into different sub-groups of magnitudes with a interval of 0.5 mag, the mean measurement errors at specified magnitudes were derived. We found that the measurement errors for each filter band are larger at fainter depth. Typically, the errors are $0.02^m$ for bright stars (say, $m<16.5^m$), and $0.05^m$ at $m=19.0^m$.  For magnitudes in $c$ and $o$ filters, the errors from the model calibration (Zhou et al. 1999) should be taken into account. By comparing the observed SEDs and the template SEDs, the typical errors are $0.05^m$ and $0.10^m$ respectively for $c$ and $o$ filters, which should be larger than the upper limits of the observational errors because the difference between observed SEDs and template SEDs includes observational measurements errors, template SED errors, and data fitting errors. The principle of this kind of error estimation can be found in Zhou et al. (2001). ",
        "conclusions": "This paper presents our multicolor optical photometry for the central region ($56' \\times 56'$) of Abell 2634, using the 60/90 cm Schmidt Telescope of Beijing Astronomical Observatory equipped with 14 BATC filters which cover almost the whole optical wavelength domain. As a result, we obtained the SEDs of 5572 objects. With the help of NASA/IPAC Extragalactic Database and the compilations in the literature, we selected 124 known member galaxies from our SED lists, for which the detailed analyses on their SED features are performed. The early-type galaxies are dominant in the central region of Abell 2634. The sample of 59 known early-types offers an opportunity to check the accuracy and reliability of our SED data by the comparison between the observed SEDs and template SEDs. We applied the SED fitting technique on 124 known member galaxies to estimate the photometric redshifts, and obtained good results: 80\\% of the members have the $z_{phot}$ less than 0.06, which provides a reliable tool for further membership selection. Furthermore, a verification of color-magnitude correlation using our SED information of early-types demonstrates the great advantages and potential for BATC multicolor observations in studying the color-magnitude effect.  Based on the knowledge of SED features of known member galaxies, we carefully utilized the color-color diagrams, technique of photometric redshift, and color-magnitude correlation (particularly for early-type members) in the selection of faint members. As a result, we isolated 74 galaxies as the faint members, including 58 early-types and 16 spirals. Then, based on the enlarged sample of member galaxies, the spatial distributions for morphologically various galaxies are discussed. We found that most of blue galaxies are scattered in  the outer region of this cluster, and core region is populated by early-type  galaxies for Abell 2634. Furthermore, the color dispersion for early-type galaxies is larger in the outer region, which might be a kind of environmental effect on the evolution of cluster galaxies.  Based on our SED catalogs of known and new members, some studies on the evolution of galaxy populations in the Abell 2634 can be carried out. In order to investigate the dynamics and kinematics of the galaxies in the central region, the follow-up efforts of spectroscopic observations for sake of obtaining the accurate radial velocities for all the faint member galaxies can be proposed.  \\vskip 1cm"
    },
    "0201/astro-ph0201334_arXiv.txt": {
        "abstract": "It has been suggested that the anomalous Na, Mg and Al observed in Globular Cluster Red Giant stars could be the result of a thermally unstable hydrogen shell. Currently accepted reaction rates indicate that temperatures of approximately 70-75 million K are required to produce the observed enhancements in Na and Al along with depletions in Mg.  The work presented here attempts to model the H shell instability by a simple mechanism of altering the energy production in the region of the H shell. We show that even extreme cases only give rise to small intermittent temperature increases that have minimal affect on the surface abundances. Full evolutionary modelling incorporating this technique simply accelerates the evolution of the RGB phase producing the same surface abundances as other models but at an earlier time.  We conclude that, unless hydrogen shell instabilities manifest themselves quite differently, they are unlikely to lead to the required temperatures and alternative explanations of the abundance anomalies are more promising. ",
        "introduction": "As long ago as 1947, Popper \\shortcite{Popper:47a} had noted that the globular cluster red giant L199 in M13 was CN-strong. Further work by Harding \\shortcite{Harding:62a}, Osborn \\shortcite{Osborn:71a} and Hesser, Hartwick and McClure \\shortcite{Hesser:76a} confirmed that this anomaly is common. So for over fifty years it has been known that the surface abundances of elements such as C and N can vary from one red giant star to another within an individual globular cluster.  It has become clear in the intervening years, however, that these anomalies are not restricted to C and N:~heavier elements such as O, Na, Mg and Al also show star to star variations. Most other elements studied show little variation, and the degree of variation of C, N, O etc. differs from cluster to cluster and is much less apparent in field stars \\cite{Langer:92a}.  Observations by many groups, including Kraft et al. \\shortcite{Kraft:97a}, Smith and Kraft \\shortcite{Smith:96a}, Langer \\shortcite{Langer:92a}, Shetrone \\shortcite{Shetrone:96a} \\shortcite{Shetrone:96b} \\shortcite{Shetrone:97a}, have identified this abundance anomaly problem and the most commonly suggested explanation for such high Na and Al are the deep mixing \\cite{Sweigart:79a} and primordial enhancement \\cite{Cottrell:81a} hypotheses. The deep mixing hypothesis suggests that the rotation of the star can give rise to meridional circulation currents that in turn create a mixing zone in the radiative region separating the top of the hydrogen burning shell (HBS) and the base of the convective envelope (BCE). This would allow elements processed in the HBS to be mixed to the surface as the star proceeds up the RGB. The primordial enhancement hypothesis suggests that the stars are born with these anomalies already present. The stars are formed from material ejected by intermediate mass ($\\simeq 3-10 \\msun$) AGB stars from an earlier epoch. Without going into the full mechanisms of each process, either could in principle explain the observed anomalies, but modelling and observations have indicated that probably both processes occur to some extent.  Langer, Hoffman and Zaidins \\shortcite{Langer:97a} suggested that the Na and Al abundance anomaly problem may be solved if the H shell reached higher temperatures for at least part of its lifetime whilst on the RGB. The idea stems from work by Von Rudloff and Vandenburg \\shortcite{VonRudloff:88a} and even earlier work by Bolton and Eggleton \\shortcite{Bolton:73a}. Von Rudloff and Vandenburg \\shortcite{VonRudloff:88a} found that the H shell of a non-rotating small mass red giant model was stable but only marginally. A rotating model could be unstable. The link with rotation is particularly encouraging because this mechanism could explain both the deep mixing and the thermally unstable shell. The work performed by Langer et al \\shortcite{Langer:97a} shows that a thermally unstable H shell would allow for the greater production of $^{27}$Al at the expense of $^{24}$Mg and that there would also be extra $^{23}$Na available for mixing to the surface.  The conclusion was that if the H shell reached temperatures of approximately 70-75 million K then the material mixed to the surface would produce a composition in close agreement with observation. In fact temperatures of this order and no lower would be required to match the Na observations because at the lower temperatures too much $^{23}$Na would be produced. Also it is only at these temperatures that significant $^{27}$Al is produced and $^{24}$Mg is only destroyed at temperatures above 70 million K. Therefore, to have any surface depletions in $^{24}$Mg the temperature of the H shell has to reach at least 70 million K unless primordial changes are invoked or the accepted reaction rates are altered. The preferred temperature range was 72-73 million K (see section 2 of Langer et. al \\shortcite{Langer:97a}). This is the temperature at which they believed all the observed surface abundance anomalies could be matched.  As the structure of a model is heavily dependent upon the temperature profile any change to the temperature must be included in the evolution modelling. The following sections discuss models where temperatures in the H shell are artificially increased by decreasing the energy generation rate. The next section looks at oscillating temperature rises in the HBS; i.e. a pulsating model where each pulse immediately follows the previous pulse. In the following section pulses in the temperature of the HBS are created with a time gap between each pulse. ",
        "conclusions": "It appears that a thermally unstable H shell, modelled with the methods of this paper,  cannot produce the high temperatures of about 70 million K required to produce the observed aluminium abundances. We succeeded only in accelerating the evolution with negligible effect on the interior temperatures and the surface abundances. Although a genuine instability, may manifest itself in a way which is quite different to these exploratory calculations, we have tried to artificially reproduce the most favourable conditions without success.  It appears that primordial abundances variations among the stars combined with deep mixing currently offer greater potential in making progress on this problem."
    },
    "0201/astro-ph0201428_arXiv.txt": {
        "abstract": "The principles underlying a proposed class of black hole accretion models are examined.  The flows are generally referred to as ``convection-dominated,'' and are characterized by inward transport of angular momentum by thermal convection and outward viscous transport, vanishing mass accretion, and vanishing local energy dissipation.  In this paper we examine the viability of these ideas by explicitly calculating the leading order angular momentum transport of axisymmetric modes in magnetized, differentially rotating, stratified flows.  The modes are destabilized by the generalized magnetorotational instability, including the effects of angular velocity and entropy gradients.  It is explicitly shown that modes that would be stable in the absence of a destabilizing entropy gradient transport angular momentum outwards.  There are no inward transporting modes at all unless the magnitude of the (imaginary) \\BV frequency is comparable to the epicyclic frequency, a condition requiring substantial levels of dissipation.  When inward transporting modes do exist, they appear at long wavelengths, unencumbered by magnetic tension.  Moreover, very general thermodynamic principles prohibit the complete recovery of irreversible dissipative energy losses, a central feature of convection-dominated models.  Dissipationless flow is incompatible with the increasing inward entropy gradient needed for the existence of inward transporting modes.  Indeed, under steady conditions, dissipation of the free energy of differential rotation inevitably requires outward angular momentum transport.  Our results are in good agreement with global MHD simulations, which find significant levels of outward transport and energy dissipation, whether or not destabilizing entropy gradients are present. ",
        "introduction": "Originally developed to be powerful luminosity sources, black hole accretion models are now confronted by an embarrassing plethora of underluminous X-ray emitters, the best known of which is the Galactic center source Sgr A* (Melia \\& Falcke 2001).  These low luminosity objects are thought to be prime candidates for a class of theoretical accretion models that has been intensively studied in recent years, which we shall refer to generically as nonradiative accretion flows. Accretion generally requires significant energy loss, and the absence of radiative losses in these flows means that determining the ultimate fate of the gas is less than straightforward.  Much of the recent interest in nonradiative flows was sparked by the work of Narayan \\& Yi (1994), who examined a series of one-dimensional, self-similar, steady accretion models referred to as ``advection-dominated accretion flows,'' or ADAFs for short.  In these models an $\\alpha$ viscosity allows angular momentum transport and the resulting flows are quasi-spherical and substantially sub-Keplerian. Dissipative heating increases toward the flow center, creating an inwardly increasing entropy profile which we shall refer to as ``adverse.''  Nonradiative accretion flows are amenable to numerical simulation. Hydrodynamical simulations carried out with large assumed $\\alpha$ values ($>0.3$) show some similarities to ADAFs (\\cite{ia99}, 2000), but smaller values of $\\alpha$ led to flows with substantial turbulence, smaller than anticipated net inward mass accretion rates, and density distributions that are far less centrally peaked than in ADAFs (\\cite{spb99}; \\cite{ia99}, 2000).  Both inward and outward mass fluxes were observed at different times and different locations within the flow, and they nearly cancelled.  These findings were given the following interpretation by Narayan, Igumenschev, \\& Abramowicz (2000; hereafter NIA), Quataert \\& Gruzinov (2000; hereafter QG), and Abramowicz et al.~(2002; hereafter AIQN). The adverse entropy gradient triggers an instability, and significant levels of convection result, hence the global solutions were called ``convection-dominated accretion flows''  (CDAFs).  The next step in the argument is key:  invoking the findings of other hydrodynamical simulations (Stone \\& Balbus 1996; Igumenschev, Abramowicz, \\& Narayan 2000), the angular momentum transport generated by the convective turbulence was claimed to be {\\em inward}.  This inward transport by convection was envisioned to be sufficiently great as to cancel the primary outward angular momentum transport by whatever process the $\\alpha$ viscosity was modeling---presumably the magnetorotational instability, or MRI (Balbus \\& Hawley 1991).  In the CDAF scenario the vanishing of the angular momentum flux implies that the $R\\phi$ component of the stress tensor responsible for accretion also vanishes (NIA, AIQN).  This has the further consequence that there is essentially no mass accretion and no dissipation, despite the presence of vigorous turbulence throughout the bulk of the flow. The only region where there is any mass accretion in the model is at the very inner edge of the flow.  All of the energy release associated with this small net accretion is transported outward to infinity by the surrounding convective flow which is maintained with no further dissipative losses.  For this reason, CDAFs are put forth as natural candidates to explain under-luminous X-ray sources.  These models have been elaborated upon, becoming influential and widely-cited.  Since black hole accretion models are central to our understanding of much of X-ray astronomy, the theoretical foundations for CDAFs deserve careful scrutiny.  In this work we carry out an explicit analysis of magnetized, rotationally-supported gas in the presence of an adverse entropy gradient.  We find that CDAF models have two major inconsistencies. First, locally unstable disturbances with adverse entropy gradients do not generally transport angular momentum inwards in magnetized fluids. Rather, they generally transport angular momentum outwards. Qualitatively, their  behavior is indistinguishable from standard MRI modes.  This latter point has been emphasized elsewhere (Hawley, Balbus, \\& Stone 2001; Balbus 2001), but here we demonstrate it quantitatively by explicitly calculating the leading order angular momentum transport associated with unstable WKB modes.  Convective modes transport angular momentum outwards when magnetic tension is significant, and inwards only for the very longest wavelength (global scale) disturbances, where magnetic tension forces are negligible. Indeed, for a given wavenumber, the direction of angular momentum transport is less a matter of whether it is destabilized by convection or rotation, and more a matter of the nature of background medium:  is it effectively magnetized or not?  This is the crucial issue.  The second difficulty is more direct and fundamental, affecting magnetic and nonmagnetic models alike.  By relying upon dissipated heat energy to trigger a convective instability that supposedly renders the flow dissipation-free, CDAFs run afoul of thermodynamic principles.  If the source of the free energy is differential rotation, the direction of angular momentum transport {\\em must} be outward.  This is a serious inconsistency.  The dissipation is quite significant, and is indeed essential if convectively unstable entropy profiles are to be sustained.  We are led to a much more standard picture of the dynamics of turbulent accretion flows, though one at odds with the tenets of CDAF theory. The turbulent stress tensor in magnetized differentially rotating gas does {\\em not} vanish.  There is vigorous local turbulent dissipation. There is mass accretion.  The near cancellation of instantaneous inward and outward mass fluxes is a property of any turbulent flow with large rms fluctuations, and not a superposition of the contributions from two distinct sources of mass flux with opposite signs.  In the following sections, we present (\\S 2) the details of the angular momentum calculation showing outward transport; (\\S 3) an explanation of important thermodynamic inconsistencies evident in CDAF theory; (\\S 4) a brief review of numerical simulations and a concluding summary.  \\section {Radial Angular Momentum Transport}  \\subsection {Local WKB Modes} Consider a disk with radially decreasing outward entropy and pressure gradients.  We use standard cylindrical coordinates $(R, \\phi, Z)$.  The square of the \\BV frequency ($N^2$) is thus negative, and tends to destabilize.  In what follows, it is convenient to work with the positive quantity \\beq {\\cal N}^2 \\equiv -N^2  = {3\\over 5\\rho}{\\dd P\\over\\dd R}{ \\dd\\ln P\\rho^{-5/3}\\over\\dd R}> 0. \\eeq  The disk is differentially rotating with decreasing outward angular velocity $\\Omega(R)$, and epicyclic frequency \\beq \\kappa^2 = 4\\Omega^2 + {d\\Omega^2\\over d\\ln R} = {1\\over R^3} {dR^4\\Omega^2\\over dR} > 0. \\eeq A vertical magnetic field $\\bb{B} = B \\bb{e_Z}$ threads the disk. Its associated Alfv\\'en velocity is  $v_A^2 = {B^2/4\\pi\\rho}$, where $\\rho$ is the gas density. Axisymmetric WKB plane wave displacements of the form \\beq \\bb{\\xi} (R, Z, t) = \\bb{\\xi} \\exp (i{kZ} - i \\omega t), \\eeq where ${k}$ and $\\omega$ are respectively the vertical wavenumber vector and the angular frequency, lead to the dispersion relation (Balbus \\& Hawley 1991) \\beq\\label{disp} {\\tilde\\omega}^4 +{\\tilde\\omega}^2({\\cal N}^2 -\\kappa^2) -4\\Omega^2 (kv_A)^2 = 0 , \\eeq where \\beq {\\tilde\\omega}^2 =\\omega^2 -(kv_A)^2. \\eeq  Let $\\gamma= - i\\omega$. Then, the unstable branch of the dispersion relation (\\ref{disp}) is \\beq\\label{growth} \\gamma^2 = -(kv_A)^2 +{1\\over2} \\left[ {\\cal N}^2 -\\kappa^2 + \\sqrt{({\\cal N}^2 - \\kappa^2)^2 +16\\Omega^2(kv_A)^2}\\right]. \\eeq It is straightforward to show that these unstable modes must have \\beq (kv_A)^2 < {\\cal N}^2 - {d\\Omega^2\\over d\\ln R} = {\\cal N}^2 + \\left|{d\\Omega^2\\over d\\ln R}\\right| , \\eeq and that the maximum growth rate is \\beq \\gamma_{max} = {\\Omega\\over 4} \\left( {{\\cal N}^2\\over\\Omega^2 } + \\left|{d\\ln\\Omega^2\\over d\\ln R}\\right| \\right), \\eeq which is attained for wavenumbers satisfying \\beq\\label{mg} (kv_A)_{max}^2 = \\Omega^2 \\left( 1 - {( {\\cal N}^2 -\\kappa^2)^2\\over 16\\Omega^4} \\right). \\eeq  \\subsection{Stress Calculation}  The angular momentum flux is directly related to the $R\\phi$ component of the stress tensor \\beq T_{R\\phi}  = \\rho ( \\delta v_R\\, \\delta v_\\phi - \\delta v_{AR}\\, \\delta v_{A\\phi}), \\eeq where $\\delta$ denotes an Eulerian perturbation, and \\beq \\bb{\\delta v_A} = { \\bb{\\delta B}\\over \\sqrt{4\\pi\\rho}}. \\eeq For the local WKB modes we consider here, the angular momentum flux is $R\\Omega T_{R\\phi}$.  Hence, the sign of the transport is simply the sign of $T_{R\\phi}$.  The needed expressions can be written down immediately from the equations (2.3c--g) of Balbus \\& Hawley (1991).  In terms of $\\gamma$, they are \\beq\\label{Rey} \\delta v_\\phi \\delta v_R = (\\delta v_R^2) {\\Omega\\over D \\gamma} \\left( {(kv_A)^2\\over\\gamma^2} \\left|{d\\ln \\Omega\\over d\\ln R}\\right| - {\\kappa^2\\over 2\\Omega^2} \\right), \\eeq \\beq\\label{Max} -\\delta v_{A\\phi} \\delta v_{AR} = (\\delta v_{AR})^2 {2\\Omega\\over D \\gamma} = (\\delta v_{R})^2 \\left(kv_A\\over \\gamma\\right)^2  {2\\Omega\\over D \\gamma} , \\eeq where \\beq D= 1 + {(kv_A)^2\\over \\gamma^2}. \\eeq These equations are general beyond our simple example, holding in the presence of both vertical and radial entropy gradients.  Note that there is no explicit dependence upon ${\\cal N}$; the only dependence upon ${\\cal N}$ anywhere is through the growth rate $\\gamma$.  The Maxwell stress (\\ref{Max}) must always be positive.  For a given growth rate, angular momentum transport is completely determined by rotation and magnetic tension.  Figure (1) shows the stability and angular momentum transport properties of an unmagnetized Keplerian disk in the $k^2-{\\cal N}^2$ plane.  The vertical scale is immaterial: when ${\\cal N}>\\kappa$, convectively unstable modes are triggered at all wavenumbers.  There is nothing special about the physics of convection {\\em per se} with regard to inward angular momentum transport.  In the absence of a magnetic field, any axisymmetric disturbance governed by the above equations would transport angular momentum inwards.  This is why hydrodynamic simulations consistently find inward transport.  \\begin{figure} \\epsscale {1.05} \\plotone{hydro.eps} \\caption{Region of instability in the $k^2$ --  $|N^2|/\\Omega^2$ plane, for an unmagnetized disk with a Keplerian rotation profile.  The system is unstable for all wavenumbers when $|N^2| > \\kappa^2$, and the sense of transport is inward.}\\label{stability} \\end{figure}  \\begin{figure} \\epsscale {1.05} \\plotone{mhd.eps} \\caption{ Region of instability in the $(kv_A/\\Omega)^2$ -- $|N^2|/\\Omega^2$ plane, for a magnetized disk obeying a Keplerian rotation law.  (Other rotation profiles lead to qualitatively similar diagrams.) The hatched region indicates the domain of outward transport.  The solid line labeled $\\gamma_{max}$ shows the wavenumber of maximum growth rate as a function of ${\\cal N}^2$.  The thin wedge at the bottom corresponds to the inward transport region.  See the text for the definitions of regions A through D.  } \\label{mstability} \\end{figure}    Matters change radically when a magnetic field is present. By combining equations (\\ref{growth}), (\\ref{Rey}), and (\\ref{Max}), we may calculate the full range of wavenumbers that transport angular momentum outwards: \\beq\\label{out} (kv_A)^2 > {1\\over 16}\\left(4-\\left|d\\ln\\Omega^2\\over d\\ln R \\right|\\right) \\left(2{\\cal N}^2 - \\kappa^2\\right). \\eeq It is apparent that unless ${\\cal N}$ is sustained at values in excess of $\\kappa/\\sqrt{2}$, every mode, convective or otherwise, will transport angular momentum outward.  The content of equation (\\ref{out}) is shown graphically in figure (2). Clearly, the dominant direction of angular momentum transport is now outwards.  Note that the sense of transport is outward even for those wavenumbers that satisfy the condition \\beq\\label{conmd} \\left|{d\\Omega^2\\over d\\ln R}\\right| < (kv_A)^2 < {\\cal N}^2 + \\left| {d\\Omega^2\\over d\\ln R}\\right| \\eeq which are unstable {\\em only} because of the presence of an adverse entropy gradient.  These wavelengths occupy Region A in the figure, specifically the shaded wedge above the dotted line $(kv_A)^2 = 3\\Omega^2$.  These wavelengths would otherwise be stable to the ``pure MRI.'' They are destabilized only because of the existence of adverse entropy gradients, yet all region A modes transport angular momentum outwards.  Physically this is because the large magnetic tension, which regulates angular momentum transport in this regime, would oridinarily be a strongly stabilizing agent.   The adverse entropy gradient destabilizes, but does not alter the outward flow of angular momentum.  Region B denotes the shaded area below region A, but above the unshaded wedge comprising regions C and D.  The B region is the wavenumber domain of outward transporting modes that would be unstable in the presence of the pure MRI.  Their growth rate is increased by the presence of finite ${\\cal N}^2$.  The curve showing the most rapidly growing wavenumber (equation [\\ref{mg}]) lies in region B, as shown, but eventually leaves this region at larger values of ${\\cal N}^2$ (see below).  The inward transporting modes are confined to the unshaded narrow wedge at the bottom of the diagram, corresponding to long wavelengths.  These modes have only a small magnetic tension force, $(kv_A)^2 \\ll \\Omega^2$.  Region C consists of modes with ${\\cal N}^2 < \\kappa^2$. Comparison with figure (1) shows that region C modes are destabilized only by the MRI (i.e., they are stable in hydro disks), but are nevertheless associated with inward transport.  Finally, region D identifies the modes that would be unstable in a purely hydrodynamic system that likewise transport angular momentum inward.  In this narrow domain, both magnetic tension and rotational stabilization are smaller than the adverse entropy gradient.  The curve denoting the most rapidly growing wavenumber enters region D for values of ${\\cal N}^2$ in excess of $4\\Omega^2$.  The region of inward transport is very small, indeed non-existent for ${\\cal N}^2<0.5 \\kappa^2$, and it would be surprising if these marginally-valid WKB modes effectively halted angular momentum transport in nonradiative flows.  In the next section we show that this is in fact impossible in any system where the seat of free energy is differential rotation.  For the present, it is useful to have an estimate (or a bound) of the size of ${\\cal N}^2$ one expects in nonradiative flows.  The entropy equation for a monotomic gas is \\beq {3\\over2} P {d\\ \\over dt} \\left(\\ln P\\rho^{-5/3}\\right) =  Q^{+} \\eeq where $Q^+$ represents dissipative heating.  This will generally not exceed the total energy budget available from differential rotation, $-T_{R\\phi} d\\Omega/d\\ln R$, and may be much less. In a one-dimensional approximation, we therefore expect \\beq {3\\over2} P v_r {d \\over dr} \\left(\\ln P\\rho^{-5/3} \\right)  \\lta -T_{R\\phi} {d\\Omega\\over d\\ln R} \\eeq Let us work near the equatorial plane and switch to cylindrical radius $R$. The inward drift velocity is related to the stress tensor by (Balbus \\& Hawley 1998): \\beq v_R \\simeq - {T_{R\\phi}\\over \\rho R \\Omega}, \\eeq which leads to \\beq {1\\over\\rho} {d \\over dR} \\left(\\ln P\\rho^{-5/3} \\right) \\lta {1\\over3} {R\\over P} {d\\Omega^2 \\over d \\ln R}, \\eeq and \\beq {\\cal N}^2 = {3\\over 5\\rho} {dP\\over dR} {d \\over dR} \\left(\\ln P\\rho^{-5/3} \\right) \\lta {1\\over 5} {d\\ln P\\over d \\ln R} {d\\Omega^2\\over dR}. \\eeq Generally, when differential rotation increases, the pressure gradient decreases, and vice-versa.  For a Keplerian profile, ${\\cal N}^2/\\Omega^2 \\lta 0.6 d\\ln P/d\\ln R$.  In such a disk, pressure gradients are small, and ${\\cal N}^2/\\Omega^2$ is likely to be less than unity.  If a significant fraction of free energy goes into generating a magnetic field that is carried off from the disk (i.e., not into dissipative field reconnection), then ${\\cal N}^2$ could be appreciably smaller. The point here is that a well-defined bound on the rate of energy dissipation limits the size of ${\\cal N}^2$. ",
        "conclusions": ""
    },
    "0201/hep-ph0201178_arXiv.txt": {
        "abstract": "The past few years have seen dramatic breakthroughs and spectacular and puzzling discoveries in astrophysics and cosmology.  In many cases, the new observations can only be explained with the introduction of new fundamental physics.  Here we summarize some of these recent advances.  We then describe several problem in astrophysics and cosmology, ripe for major advances, whose resolution will likely require new physics. ",
        "introduction": "The goal of the Snowmass 2001 P4 working group was to identify opportunities for advances at the interface of particle physics, astrophysics, and cosmology. Since the previous Snowmass meeting (Snowmass96 \\cite{sm96}), there have been spectacular advances in cosmology and particle astrophysics.  These include, but are not limited to, cosmic microwave background (CMB) evidence favoring inflation, supernova and CMB evidence for negative-pressure dark energy, and results from solar- and atmospheric-neutrino experiments.  Taken together, these results demonstrate that astro/cosmo/particle physics is an integral component of the particle-physics research enterprise.  The P4 Working Group covered a very broad range of topics, subdivided into eight topical groups: \\begin{enumerate} \\item Dark matter and relic particles \\item Gamma rays and X-rays \\item Cosmic microwave background and inflation \\item Structure formation and cosmological parameters \\item Cosmic rays \\item Gravitational radiation \\item Neutrino astrophysics \\item Early Universe and tests of fundamental physics \\end{enumerate} There is intense experimental and theoretical activity in all of these areas.  Every major breakthrough listed below has been made since Snowmass96, and the prospects for the next decade and beyond are even brighter.   It is not possible to describe here all the work discussed during the P4 sessions at Snowmass. Instead, we emphasize the overarching themes that these areas represent and their relevance to the purpose of Snowmass 2001.  There have been spectacular observational breakthroughs:  \\begin{itemize} \\item Recent CMB measurements provide evidence that the total energy density of the Universe, $\\Omega_{\\rm tot}$, is close to unity.  For the first time, we may know the geometry of the Universe.  Observations support the hypothesis that large-scale structure grew from primordial density fluctuations, in agreement with predictions from inflation.  This provides the scientific connection between the large-scale structure of the Universe and elementary particle physics, and is indicative of new physics at higher energy scales.  \\item The discrepancy between a matter density $\\Omega_m \\simeq 0.3$ and $\\Omega_{\\rm tot} \\simeq 1$ provides independent corroboration of the remarkable recent supernova-survey evidence for some form of ``dark energy''.  If confirmed, this suggests that $\\sim70\\%$ of the energy density of the Universe is of a previously unknown and mysterious type.  The existence of the dark energy was not even suspected by most physicists at the time of Snowmass96.  \\item The CMB data verify that 25\\% of the density of the Universe must be in the form of nonbaryonic dark matter---as suggested earlier by dynamical measurements of the matter density and big-bang-nucleosynthesis predictions of the baryon density---implying physics beyond the standard model.  This strengthens the case for some form of particle dark matter (e.g., supersymmetric particles or axions) in our Galactic halo.  \\item At the same time that the case for nonbaryonic dark matter continues to strengthen, prospects for detecting dark-matter particles over the next decade, using both direct and indirect methods, are promising.  A fleet of experiments with complementary sensitivities and systematics are probing deeper into theoretically well-motivated regimes of particle-physics parameter space.  \\item Underground observations of solar neutrinos and cosmic-ray-induced atmospheric neutrinos, indicating the existence of neutrino oscillations, have provided evidence that neutrinos are not massless.  This is the first direct experimental confirmation that the standard model is incomplete.  \\item The most massive black holes are central to TeV-class astrophysical accelerator systems that have been observed to radiate immense power in gamma rays.  Unprecedented leaps in sensitivity will be made by the next generation of both ground-based and space-based gamma-ray instruments, opening up large discovery spaces.  Together, these gamma-ray measurements also provide a unique probe of the era of galaxy formation, and will provide the first significant information about the high-energy behavior of gamma-ray bursts.  \\item The evidence for the highest-energy cosmic-ray events ($E>10^{20}$ eV) poses significant challenges to our theoretical understanding.  Observations under way, and new detectors under construction and in planning, both on the ground and in space, will shed new light on this highest-energy mystery.  These experiments, as well as high-energy neutrino telescopes, will allow us to exploit the highest-energy particles for particle physics.  \\item Soon, the Universe will be viewed not only with photons, but also with high-energy charged particles, high-energy neutrinos, and gravitational waves.  These observations will test strong-field general relativity and open vast new windows on the highest-energy astrophysical phenomena and the early Universe. \\end{itemize}  Interest in astro/cosmo/particle physics has grown remarkably since Snowass96.   Of the five P working groups at Snowmass 2001, the subscription to P4 was second only to that of P1, Electroweak Symmetry Breaking, and the overlap with all of HEP was obvious in P4 joint sessions with other groups.  Interest in the Snowmass-wide teach-in on astro/cosmo/particle physics was very strong.  In summary, there is vigorous and fast-growing activity in astro/cosmo/particle physics, and the intellectual overlap between the particle-physics community and others---especially the high-energy astrophysics and cosmology communities---has grown into full and healthy partnerships, greatly accelerating progress.  These partnerships provide enormous opportunities as well as new challenges. ",
        "conclusions": "Particle astrophysics and cosmology present a number of exciting theoretical and experimental challenges along a broad front. There are a variety of mysterious astrophysical and cosmological phenomena that almost certainly will require new physics to be understood. Here we have surveyed much of the activity in this field, focusing in particular on areas with recent breakthroughs, and presented a broad overview of some of the most exciting questions that may be addressed experimentally in the foreseeable future.  More detailed discussions of all of the topics we have covered can be found in the topical subgroup reports in this Snowmass proceedings."
    },
    "0201/astro-ph0201217_arXiv.txt": {
        "abstract": "The recent availability of special purpose computers designed for calculating gravitational interactions of $N$-bodies at extremely high speed has provided the means to model globular clusters on a star-by-star basis for the first time. By endeavouring to make the $N$-body codes that operate on these machines as realistic as possible, the addition of stellar evolution being one example, much is being learnt about the interaction between the star cluster itself and the stars it contains. A fascinating aspect of this research is the ability to follow the orbits of individual stars in detail and to document the formation of observed exotic systems. This has revealed that many stars within a star cluster lead wildly promiscuous lives, interacting, often intimately and in rapid succession, with a variety of neighbours. ",
        "introduction": "\\label{s:intro}  The rich environment of a star cluster provides an ideal laboratory for the study of self-gravitating systems. Star clusters within our Galaxy range in size from loose associations and open clusters, which contain up to tens of thousands of stars, to the large globular clusters which host a million stars or more. Globular clusters are some of the oldest objects known and dynamically they lie in a very interesting regime. Compared to the solar neighbourhood the central density of stars is high enough (a factor of 10 million or more greater in some globular clusters) that a significant fraction of the cluster stars are likely to experience at least one close encounter with another star during their lifetime. At the opposite end of the scale, the size of a globular cluster is small enough, in comparison to a galaxy, that its dynamical relaxation timescale is less than its age. These combined considerations make a globular cluster an exciting place for a star to reside. As such we expect the stellar populations of a star cluster, the various classes of stars and binaries, to exhibit a dynamical signature. In particular, many of the stars contained in the cluster will deviate strongly from the evolutionary paths predicted by standard stellar and binary evolution theory. This is indeed what we observe: either when pointing the Hubble Space Telescope (HST) at the centres of globular clusters or when generating star-by-star models of star cluster evolution. ",
        "conclusions": "\\label{s:conclu}  The introduction of the GRAPE-6 hardware, coupled with $N$-body codes that have the capability to produce realistic cluster models, places the astrophysicist in the position of knowledgeable voyeur. An exciting aspect of this is the capability to investigate and understand the range of stellar populations that are observed, and how these are affected by dynamical interactions between cluster stars. As we move to full utilization of the GRAPE-6, and the simulation of globular clusters, the fascinating tales of promiscuous stars in clusters will become very graphic indeed!"
    },
    "0201/astro-ph0201351_arXiv.txt": {
        "abstract": "We present the results of a {\\it Far Ultraviolet Spectroscopic Explorer (FUSE)} survey of \\ion{O}{6} 1031.93~\\AA\\ and 1037.62~\\AA\\ absorption toward 18 OB stars in the Small Magellanic Cloud (SMC). The \\FUSE\\ data are of very high quality, allowing a detailed study of the coronal temperature gas in the SMC. We find that \\ion{O}{6} is ubiquitous in the SMC, with a detection along every sight line. The average value of the \\ion{O}{6} column density in the SMC is log $<N$(\\ion{O}{6})$>=14.53$. This value is $1.7$ times higher than the average value for the Milky Way halo (perpendicular to the Galactic plane) of log $N_\\perp$(\\ion{O}{6})=14.29 found by \\fuse, even though the SMC has much lower metallicity than the Galaxy. The column density in the SMC is higher along sight lines that lie close to star-forming regions, in particular NGC~346 in the northern part of the SMC, and to a lesser degree the southwestern complex of \\ion{H}{2} regions. This correlation with star formation suggests that local processes have an important effect on the distribution of coronal gas in the SMC. If the sight lines within NGC~346 are excluded, the mean column density for the SMC is log $N$(\\ion{O}{6})$=14.45$, only $1.4$ times higher than the Milky Way average. The standard deviation of the column densities for sight lines outside of NGC~346 is $\\pm27$\\%, somewhat lower than the deviation seen in the Milky Way halo. The lowest \\ion{O}{6} column densities, log $N$(\\ion{O}{6})$\\sim14.3$, occur in the central region and in the southeastern ``Wing'' of the galaxy. Even these low column densities are as high as the Milky Way average, establishing the presence of a substantial, extended component of coronal gas in the SMC. The \\ion{O}{6} absorption is always shifted to higher velocities than the main component of lower ionization gas traced by \\ion{Fe}{2} absorption. The \\ion{O}{6} line widths are broader than expected for pure thermal broadening at $3\\times10^5$~K, the temperature at which the \\ion{O}{6} peaks in abundance, so large non-thermal motions or multiple hot gas components are likely present. We discuss several mechanisms that may be able to explain the observed properties of the hot gas, including supershells, a galactic fountain, and the infall of gas previously stripped from the SMC by tidal interactions with the Milky Way and the Large Magellanic Cloud. If a galactic fountain produces the hot gas, the mass flux per unit surface area is $\\dot M/\\Omega \\sim 2\\times10^{-2}$~M$_{\\odot}$~yr$^{-1}$~kpc$^{-2}$. ",
        "introduction": "The \\ion{O}{6} 1031.93~\\AA\\ and 1037.62~\\AA\\ lines are important diagnostics of the processes responsible for distributing mass and energy in the interstellar medium (ISM). \\ion{O}{6} is difficult to produce through photoionization, requiring photons with $h\\nu \\ge 114$ eV. Gas containing \\ion{O}{6} is therefore most likely collisionally ionized.  Since \\ion{O}{6} peaks in abundance at $3\\times10^5$~K \\citep{sd93}, an unstable region of the cooling curve, gas containing \\ion{O}{6} is either cooling from higher temperatures, or being heated. As such, it is a useful probe of energetic processes in the ISM including shock heating \\citep{sm79}, conductive heating and cooling \\citep{bbf90}, and turbulent mixing layers \\citep{ssb93}. Coronal gas traced by \\ion{O}{6} is intermediate in temperature between the very hot ($T\\sim10^{6-7}$~K) X-ray emitting gas and the warm ($T\\sim10^4$~K) ionized medium.   Opportunities to observe the \\ion{O}{6} lines have been limited. The first observations of \\ion{O}{6} were made with the {\\it Copernicus} satellite along sight lines toward nearby bright stars ($V\\le7$, d $\\la2000$~pc) \\citep{jm74,j78a,j78b}. Later observations were possible with the ORFEUS-SPAS missions \\citep{hb96,h98,w98,ssh99}. These observations, combined with observations of \\ion{C}{4} $\\lambda\\lambda1548.20, 1550.77$, \\ion{N}{5} $\\lambda\\lambda1238.82, 1242.80$, and \\ion{Si}{4} $\\lambda\\lambda1393.76, 1402.77$, have built upon the original idea of a Galactic corona of hot gas which provides pressure confinement for high-latitude clouds \\citep{s56}. While results for \\ion{C}{4}, \\ion{N}{5}, and \\ion{Si}{4} have greatly increased our understanding of the Galactic corona, the relative lack of information on \\ion{O}{6} has been frustrating, since it probes a hotter temperature range and is less likely to be produced by photoionization.  With the launch of the {\\it Far Ultraviolet Spectroscopic Explorer (FUSE)} in 1999 \\citep{m00}, access to the \\ion{O}{6} lines has returned. \\FUSE\\ is much more sensitive than {\\it Copernicus}, allowing more distant stars and extragalactic sources to be observed. One of the major science goals of \\FUSE\\ is to study the properties and distribution of hot gas in the local universe through \\ion{O}{6} absorption. Early results have already been published for the Galactic halo and high velocity clouds \\citep{s00,sem00}, and more comprehensive surveys are in progress.   The sensitivity of \\FUSE\\ makes it possible to observe early type stars in the Magellanic Clouds, so we can now study the \\ion{O}{6} absorption in galaxies with different properties than our own. The Small Magellanic Cloud (SMC) has properties that are vastly different from the Milky Way. It has low mass, low metallicity, and it is currently undergoing a gravitational interaction with two more massive galaxies. These factors may influence the properties of hot gas in the ISM of the SMC. The hot gas in the SMC has been studied previously using {\\it International Ultraviolet Explorer (IUE)} spectra \\citep{ds80,fs83,f84,f85,fs85}.  Absorption by \\ion{C}{4} and \\ion{Si}{4} was measured toward the nine stars studied by \\cite{fs85}, who found evidence for a global component of highly ionized gas. More recently, Space Telescope Imaging Spectrograph (STIS) spectra of one of the sight lines studied by \\cite{fs85}, AV229 (HD 5980), were analyzed by \\cite{k01} and \\cite{hshb01}, who also found evidence for hot gas in a supernova remnant in addition to the general ISM. The spectra show absorption by \\ion{N}{5} in addition to \\ion{C}{4} and \\ion{Si}{4}. \\fuse\\ data on \\ion{O}{6} toward Sk~108 in the SMC has already been analyzed by \\cite{m01}, but the complicated stellar continuum prevented an accurate determination of the \\ion{O}{6} column density. \\cite{w91} and \\cite{ww92} discovered an X-ray halo around the SMC using {\\it Einstein} data, with spectral properties that suggest the presence of hot, thermal ISM components.  In this paper we present the results of a \\FUSE\\ \\ion{O}{6} absorption survey of 18 stars in the SMC.  By observing stars in a variety of environments in the SMC, we search for a global component of hot gas, as well as clues to the origins of the highly ionized gas. In section 2 we describe the observations and data processing. Section 3 describes the \\ion{O}{6} content of the SMC, its variation with location, and a comparison to the Milky Way. The kinematics of the \\ion{O}{6} absorbing gas are described in section 4, and section 5 contains a discussion of the observed \\ion{O}{6} absorption. Section 6 summarizes the primary conclusions of our study. In a parallel effort, \\cite{hsfs02} present a study of \\ion{O}{6} absorption in the Large Magellanic Cloud. ",
        "conclusions": "In this paper we have described the properties of \\ion{O}{6} absorption along 18 sight lines in the Small Magellanic Cloud, using far ultraviolet spectra obtained with \\fuse. The main results are:  \\begin{enumerate}  \\item{Absorption from \\ion{O}{6} is seen along every sight line in the SMC. The sight lines range in environment from star-forming regions (\\eg, NGC~346) to the field (\\eg, AV423), and includes the periphery of the SMC (Sk~188). The occurrence of \\ion{O}{6} absorption in every direction indicates the presence of a widespread component of coronal temperature gas in the SMC. }  \\item{The average column density is log $N$(\\ion{O}{6}) $\\sim14.53$. This is about 1.7 times higher then the Milky Way average of the column density perpendicular to the plane of log $N$(\\ion{O}{6}) $\\sim14.29$ \\citep{s00}. However, if NGC~346 is excluded, the average for the rest of the SMC falls to $\\sim14.45$, only 1.4 times higher than the Galactic average. }  \\item{The column density is correlated with position in the SMC. The highest values are seen toward the star-forming region NGC~346, and to a lesser extent toward the star-forming regions in the southwestern end of the Bar. Lower values are seen in the central region and in the Wing, where there is little star formation activity. This suggests that local processes strongly affect the hot gas distribution in the SMC, particularly in NGC~346.}  \\item{The \\ion{O}{6} line widths are broader than expected for pure thermal broadening, indicating that either non-thermal motions are prevalent, or that there are multiple components of hot gas at different velocities. The SMC \\ion{O}{6} line widths are broader than the Milky Way \\ion{O}{6} lines.}  \\item{The kinematics suggest that the \\ion{O}{6} is not associated with the gas producing the bulk of the low ionization absorption at $\\sim+130$~\\kms. The \\ion{O}{6} velocity is closer that of the $\\sim+180$~\\kms\\ low ionization component, but evidence for a physical association is not conclusive. We discuss several mechanisms that may explain the observations, with superbubbles and/or a galactic fountain being the most probable.}  \\item{If a galactic fountain is responsible for producing the hot gas, the rate of mass flow per unit surface area of the flow region is $\\dot M/\\Omega \\sim 2\\times10^{-2}$ M$_{\\odot}$ yr$^{-1}$ kpc$^{-2}$, assuming an initial gas density of 10$^{-2}$ cm$^{-2}$, which is very uncertain. Since the \\ion{O}{6} column density in the SMC is within a factor of 2 of that of the Milky Way, the mass flow rates are probably similar if other conditions are the same.}   \\end{enumerate}"
    },
    "0201/astro-ph0201398_arXiv.txt": {
        "abstract": "We consider the stability of an accretion disk wind to cloud formation when subject to a central radiation force.  For a vertical launch velocity profile that is Keplerian or flatter and the presence of a significant radiation pressure, the wind flow streamlines cross in a conical layer. We argue that such regions are highly unstable, and are natural sites for supersonic turbulence and, consequently, density compressions. We suggest that combined with thermal instability these will all conspire to produce clouds. Such clouds can exist in dynamical equilibrium, constantly dissipating and reforming.  As long as there is an inner truncation radius to the wind, our model emerges with a biconical structure similar to that inferred by Elvis (2000) for the broad line region (BLR) of active galactic nuclei (AGN). Our results may also apply to other disk-wind systems. ",
        "introduction": "Non-spherical outflows are a ubiquitous feature of AGN. In this Letter we discuss the stability of such an outflow launched normally from the disk with a decaying power law velocity profile. The motivation for this work is to understand the conditions under which the flow may be unstable to cloud formation, and to infer the structure of the region in which such clouds may reside. \\citet{Elvis} has argued that data require AGN BLR clouds to reside in a narrow biconical structure. Our work herein supports this possibilty.  A variety of launching mechanisms for AGN accretion disk outflows have been studied: radiatively accelerated outflows \\citep{Arav94}, hydrodynamic line-driven winds \\citep{Proga99, Proga1}, hydromagnetic disk winds \\citep{Konigl94}, \\citep{Pudritz92}, thermal wind-type outflows in low luminosity AGNs \\citep{Pietrini}, and others. In this work we parameterize the disk wind independent of the launching mechanism.  A disk wind outflow can become unstable from linear and nonlinear perturbations in velocity, density, temperature, ionization state, etc. These can lead to cloud formation. The cloud model is a leading, though not unanimously accepted, paradigm for the structure of the AGN BLR regions: models exist with \\citep{Elvis, Urry95} and without \\citep{Murray1, Murray2} clouds. The main reason for the lack of agreement is the uncertainty about cloud survival in AGN enviroments \\citep{Mathews}. They may be unstable to evaporation (e.g., see \\citep{Pier95}), or be insufficiently supported by pressure and be dynamically unstable to shredding \\citep{Mathews, Poludnenko}. Among the suggested solutions has been magnetic confinement of clouds \\citep{Rees87, Bottorff1}.  However a key point is that a typical cloud needs only to exist long enough to reprocess radiation \\citep{Rees87, Celotti, Kuncic1, Kuncic2}. If clouds are destroyed thereafter and new clouds are formed, then the system can exist in dynamical equilibrium even though each cloud loses its identity swiftly.  This may apply to the broad line region. Short cloud survival and regeneration times can be associated with turbulence \\citep{Bottorff2, Bottorff3}. The balance between formation and destruction maintains a constant BLR cloud number density.  Our analysis suggests that a nonlinear wind instability resulting from the combination of launching profile + radiation pressure can lead to the formation of a turbulent biconical zone that might manifest itself observationally as the BLR. This biconical structure is the region where flow streamlines cross. Such a bicone is remarkably similar to that inferred empirically by \\citet{Elvis} and from numerical simulations by \\citet{Proga1}. Growing observational data corroborates the existance of such a structure (e.g. NGC 1068 \\citep{Arribas96}, Mrk 3 \\citep{Ruiz01}, etc.)  We formulate the problem in section 2.1 and describe the flow field in section 2.2. We analyze the linear and nonlinear stability of the flow and a possible scenario for thermal instability resulting in formation of the BLR clouds in section 2.3. In section 3 we discuss the results and their implications for the AGN structure and dynamics. ",
        "conclusions": "We considered the linear and nonlinear stability of a disk outflow and found that it is linearly stable to infinitesimal velocity perturbations. However, when $n>-1/2$ and a supereddington radiation source is present, nonlinear instability leads to the formation of a biconical zone of compressible turbulence. High density contrasts in this zone may trigger thermal instability and lead to the further condensation of clumps into BELR clouds. Turbulence plays a key role in establishing the dynamical equilibrium which maintains a steady cloud number density, even though each cloud is short lived.  In the regimes when the outflow can develop the turbulent bicone, the geometry is remarkably similar to that described by \\citet{Elvis}. For a disk outflow with $K \\approx 0.05$ and an initial launch velocity profile slightly flatter than Keplerian, namely $n=-0.45$, the turbulent bicone inclination angle is $\\Theta \\approx 28^o$ with the divergence angle $\\Delta\\Theta \\approx 7^o$. This is in quantitative and qualitative agreement with \\citet{Elvis}: (1) no absorbers are along the line of sight passing above the outflow; (2) along the line of sight passing through the turbulent bicone broad absorption and emission line features will be observed; (3) narrow absorption lines (NAL) will be seen along sight lines that fall inside of the bicone. Figure~\\ref{streamlines} shows the geomtry of the outflow and the emerging bicone with the angle values found here.  The outflow geometry in our picture depends only on dimensionless quantities $K$ and $n$. However, in addition to $\\Theta$ and $\\Delta\\Theta$ determined above, we can also infer the inner scale by matching observed cloud velocities. We assume the same value of $n$ and $K$ as in the previous paragraph. We take $M_{BH} =10^9\\ \\msol$ and broad absorption line velocity $v_{BAL}=10^4$ \\kms (the velocity observed when the line of sight passes directly along the bicone towards the central source). In our model $v_{BAL}$ is the terminal velocity at infinity of material launched with the largest initial velocity and potential, i.e. at the point $\\lambda_0=\\lambda^*$. Setting $\\lambda_0=\\lambda^*$ in (\\ref{vstr}) and using the expressions for $v_{\\lambda}(\\lambda^*,\\lambda)$ and $v_{z}(\\lambda^*,\\lambda)$ in the limit $\\lambda \\rightarrow \\infty$, we find \\beq \\lambda^* =\\frac{2\\alpha(1+K)}{v_{BAL}^2}. \\label{lstar} \\eeq Assuming $L=1.5L_{Edd}$ and $f_{T}=0.7$ the corresponding value of $\\lambda^* =\\lambda_0 \\approx 1000 \\ a.u.$, which matches \\citet{Elvis}. The corresponding maximum launch velocity is $\\approx 2000$ \\kms. Therefore, the observed velocity in the NAL region is $\\sim 1000$ \\kms.  In principle, the paradigm presented here applies not only to luminous AGN but to any source with a centrally symmetric potential, a sufficiently luminous central radiation source, and a disk wind, e.g. disk winds in young stellar objects.  We have not discussed the role of magnetic fields, disk rotation, metallicity, wind density fall off, or the physics of line driving. These should be considered in future work."
    },
    "0201/astro-ph0201484_arXiv.txt": {
        "abstract": "We have automated the ``Synthetic Field Method'' developed by Gonzalez et al. (1998) and used it to measure the opacity of the ISM in the Local Group dwarf galaxy Sextans A by using the changes in counts of background galaxies seen through the foreground system. The Sextans A results are consistent with the observational relation found by Cuillandre et al. (2001) between dust opacity and HI column density in the outer parts of M31. ",
        "introduction": "To observe the opacity and distribution of dust independent of theoretical models for the light or dust distribution, the dust needs to be backlighted by a more distant source. The group led by Keel and White used an occulted galaxy for this purpose and their studies (Domingue et al. (1999, 2000), Keel et al. (2001) and White et al. (1992, 2000)) showed insight into the fine structure of the dust in a few galaxies.  Gonzalez et al. (1998) developed a method to use the distribution of field galaxies which they called the ``Synthetic Field Method'' (SFM); it is not restricted to the case of nearby galaxy pairs and it can supply values for the dust opacities of galaxies in the disks of spirals and irregulars in general. ",
        "conclusions": "We conclude first that the SFM can be applied with almost the same accuracy on fields with widely different exposure times. The limitations come from background galaxy clustering and confusion, not sensitivity. Secondly, the SFM can be automated to a satisfactory degree. As the clustering is the main uncertainty, applying the SFM to larger areas with similar properties (HI column density, inside or outside the spiral arm or at a certain radius) will lower these errors and give a general picture of dust distribution in irregular and spiral galaxies.\\\\  We intend now to carry out a survey using the SFM of suitable WFPC data currently in the HST archive."
    },
    "0201/astro-ph0201437_arXiv.txt": {
        "abstract": "The well-known relationship between metallicity and luminosity in active galactic nuclei (AGNs) is addressed by introducing new metallicity measurements (based on the method of Hamann \\& Ferland, hereafter HF) for a sample of narrow-line Seyfert 1 galaxies (NLS1s). Our new results, based on a sample of 162 AGNs, including nine NLS1s, indicate that while broad-line AGNs trace a metallicity--luminosity power law with an index of $\\sim0.2$, NLS1s deviate significantly from this relationship at low luminosities. Adopting the HF method based on the \\nvciv\\ line ratio, we find that NLS1 metallicities are similar to those of some high-redshift, high-luminosity quasars. We also examined the \\nivciv\\ line ratio and compared it with  \\nvciv\\ in a sample of 30 sources including several NLS1s. We find that the two do not give a consistent answer regarding the N/C abundance ratio. This result is marginal because of the quality of the data. We suggest two alternative explanations to these results: 1) The HF metallicity--luminosity dependence is not a simple two-parameter dependence and there is an additional hidden variable in this relationship that has not yet been discovered. The additional parameter may be the accretion rate, the age of the central stellar cluster or, perhaps, something else. 2) The strong line ratios involving \\ion{N}{5}~$\\lambda1240$ suggested by HF are not adequate metallicity indicators for NLS1s and perhaps also other AGNs for reasons that are not yet fully understood. ",
        "introduction": "}  Studies of emission lines in active galactic nuclei (AGNs) indicate that metallicities are typically near the solar value in their broad- line region (BLR). Accurate determinations are difficult to obtain since the line ratios depend on unknown densities and optical depths \\cite{net90}. However, Shields (1976) showed that some BLR abundances can be derived from the observed line ratios of weak nitrogen, carbon, and oxygen lines,  almost independent of the physical properties of the gas (see Hamann \\et 2002 for more details). One such line ratio is \\nivcivA\\ with a theoretical value of $\\sim0.045$ for solar metallicity.  Others are \\niiioiii, \\oiiiciv\\, and \\niiiciii, which can determine the N/O, O/C, and N/C abundances, respectively.  This method has been used in several AGN studies (e.g., Baldwin \\& Netzer 1978; Osmer 1980; Uomoto 1984), but there are practical limitations due to the weakness of these lines.  An important development in this area was achieved by Hamann \\& Ferland (1993, hereafter HF93; see also Hamann \\& Ferland 1999), who suggested alternative abundance indicators that are somewhat model dependent but much easier to obtain observationally. In particular, they have shown that the \\ion{N}{5}~$\\lambda1240$/\\ion{C}{4}~$\\lambda1549$~and \\ion{N}{5}~$\\lambda1240$/\\ion{He}{2}~$\\lambda1640$~line ratios represent the overall BLR metallicity. They also claimed that BLR metallicity tends to grow with AGN luminosity up to $\\sim10$Z$_{\\odot}$, thus implying a metallicity--luminosity (hereafter $Z-L$) relationship, in analogy with the mass--metallicity relationship observed for elliptical galaxies. At one extreme end of this relationship one finds the luminous and high- redshift quasars as AGNs having the highest metallicities. In several cases, those quasars display relatively narrow UV emission lines (e.g., Osmer 1980; Warner \\et 2002).  The low-luminosity regime in the HF93 diagram is also occupied by narrow-line Seyfert 1 galaxies (NLS1s) that have been mostly left out from previous abundance analyses. These NLS1s are defined by their extremely narrow optical permitted emission lines (FWHM$\\ltsim2000$\\,\\kms) in comparison with normal broad-line AGNs (BLAGNs; Osterbrock \\& Pogge 1985). NLS1s show extreme AGN properties; their optical emission lines put them at one extreme end of the Boroson \\& Green (1992) primary eigenvector, and they tend to display unusual behavior in other wave bands, especially in the X-ray (e.g., Boller, Brandt, \\& Fink 1996; Leighly 1999a, 1999b).  A possible explanation for the peculiar properties of NLS1s is that they have relatively low black hole (BH) masses for their luminosities and hence a very large $L/L_{\\rm Edd}$. Since only a handful of NLS1s had their BH mass measured directly using reverberation mapping techniques (Peterson \\et 2000), this is not yet fully confirmed. Another extreme NLS1 property was recently suggested, namely, unusually high metallicities (Mathur 2000 and references therein).  However, to date, this evidence remains scarce, and no systematic abundance study has yet been carried out.  In this study we present for the first time metallicity (\\'{a} la HF93) measurements for a sample of NLS1s. In \\S~\\ref{sample_prop} we define the sample properties and data analysis. In \\S~\\ref{results} we describe our new results on the $Z-L$ relationship in AGNs and attempt to answer two related questions: (1) Do NLS1s have higher metallicities compared with BLAGNs for a given luminosity? (2) Are the higher metallicities, assumed for NLS1s related to the fundamental physical properties that drive the NLS1 phenomenon? ",
        "conclusions": "} \\subsection{New Correlations Involving \\ion{N}{5}~$\\lambda1240$ \\label{N5}}  Figure~\\ref{ratios} shows the \\nvciv\\ ({\\it top}) and \\nvheii\\ ({\\it bottom}) line ratios as a function of luminosity.  The HF93 line ratios of the entire sample were also binned in ranges of 0.5 in log $\\nu L_{\\nu}$ to minimize the effect of uneven distribution in luminosity. The average positions of all objects in each of those bins are shown as large squares with error bars in Figure~\\ref{ratios}. The error bars on the large squares represent, for each axis, the standard deviation divided by the square root of the number of objects in each bin. For each HF93 line ratio--$\\nu L_{\\nu}$ log-log diagram, we performed a linear regression analysis and calculated the Pearson and Spear-  \\scriptsize \\begin{center} {\\sc TABLE 1 \\\\ Linear Regression Parameters for N\\,{\\sc v}/C\\,{\\sc iv} versus log $\\nu L_{\\nu}$} \\vskip 4pt \\begin{tabular}{lccccc} \\hline \\hline {Data Set} & {Number of} & {Pearson} & {Spearman} & {Slope} & {Constant} \\\\ {Code$\\rm ^a$} & {Objects} & {($r$)} & {($r_s$)} & {($a$)} & {($b$)} \\\\ \\hline B           & 121 & 0.70 & 0.73 & $0.19\\pm0.02$ & $-2.86\\pm0.23$  \\\\ B$+$N       & 130 & 0.55 & 0.60 & $0.13\\pm0.02$ & $-2.11\\pm0.23$  \\\\ RQQ (B)     & 105 & 0.72 & 0.74 & $0.18\\pm0.02$ & $-2.82\\pm0.23$  \\\\ RQQ (B$+$N) & 114 & 0.56 & 0.60 & $0.13\\pm0.02$ & $-2.04\\pm0.24$  \\\\ B$+$up.lim. & 137 & 0.67 & 0.69 & $0.19\\pm0.02$ & $-2.91\\pm0.24$  \\\\ B$+$N$+$up.lim. & 146 & 0.52 & 0.57 & $0.13\\pm0.02$ & $-2.17\\pm0.24$  \\\\ \\hline \\end{tabular} \\vskip 2pt \\parbox{3.485in}{% \\small\\baselineskip 9pt \\footnotesize \\indent $\\rm ^a${Data set codes are B for BLAGNs, N for NLS1s, RQQ for radio-quiet quasars, and up.lim. for upper limits on the line ratio.} } \\end{center} \\setcounter{table}{1} \\normalsize  \\scriptsize \\begin{center} {\\sc TABLE 2 \\\\ Linear Regression Parameters for N\\,{\\sc v}/He\\,{\\sc ii} versus log $\\nu L_{\\nu}$} \\vskip 4pt \\begin{tabular}{lccccc} \\hline \\hline {Data Set} & {Number of} & {Pearson} & {Spearman} & {Slope} & {Constant} \\\\ {Code$\\rm ^a$} & {Objects} & {($r$)} & {($r_s$)} & {($a$)} & {($b$)} \\\\ \\hline B          &  98 & 0.58 & 0.60 & $0.14\\pm0.02$ & $-1.37\\pm0.27$  \\\\ B$+$N      & 107 & 0.50 & 0.53 & $0.11\\pm0.02$ & $-0.95\\pm0.24$  \\\\ RQQ (B)    &  83 & 0.61 & 0.64 & $0.14\\pm0.02$ & $-1.25\\pm0.26$  \\\\ RQQ (B$+$N)&  92 & 0.54 & 0.57 & $0.11\\pm0.02$ & $-0.87\\pm0.23$  \\\\ B$+$up.lim.& 110 & 0.54 & 0.55 & $0.15\\pm0.02$ & $-1.45\\pm0.28$  \\\\ B$+$N$+$up.lim.& 119 & 0.46 & 0.48 & $0.11\\pm0.02$ & $-1.00\\pm0.26$  \\\\ \\hline \\end{tabular} \\vskip 2pt \\parbox{3.485in}{% \\small\\baselineskip 9pt \\footnotesize \\indent $\\rm ^a${Data set codes are identical to those in Table 1.} } \\end{center} \\setcounter{table}{2} \\normalsize  \\noindent man linear correlation coefficients. This analysis was carried out once for the BLAGN population and a second time for the entire data set. The results are presented in Tables 1 and 2; one can see that the exclusion of radio-loud objects or the addition of flux ratio upper limits (that were treated as real ratios) had little effect on the correlations and slopes.  Inspection of Figure~\\ref{ratios} and Tables 1 and 2 shows that the \\nvciv\\, ratio is strongly correlated with luminosity in the BLAGN case. However, it is also apparent that this relationship {\\it breaks down completely} at low luminosities, when the NLS1s are introduced into the sample. Several NLS1s that belong in the lowest luminosity regime even have \\nvciv\\, ratios as high as those of the most luminous high-z quasars. The extremely strong \\ion{N}{5}~$\\lambda1240$ in some NLS1s has been noted earlier by Wills et al. (1999) who did not investigate the resulting N/C abundance ratio. The \\nvheii\\, ratio behaves somewhat differently when NLS1s are added (Fig.~\\ref{ratios}). We find that the \\nvheii\\, ratio in BLAGNs is correlated with luminosity, although not as strong as in the \\nvciv\\, case, and that NLS1s only slightly deviate from the \\nvheii\\, $Z-L$ slope. We also discover that both line ratios are not correlated with luminosity for $\\nu L_{\\nu}$ at $1450$\\AA\\,$\\ltsim10^{46}$ erg s$^{-1}$, neither for BLAGNs nor for the entire sample. We emphasize once more that the \\nvciv\\ ratio in NLS1s is very similar to the one observed in high-$L$ AGNs and this result, which is based on an easy-to-measure line ratio, is enough to completely change the original HF93 correlation, regardless of the exact slope or value of the correlation coefficient.  \\subsection{High \\nvciv\\ at Low Luminosity \\label{reliable}}  Our new results point at two distinct scenarios: either (1) the HF93 line ratios overpredict N/C at least under some physical conditions or (2) high metallicities at low luminosities are possible and are seen in NLS1s. We discuss briefly the implications of these two scenarios and defer the more detailed analysis to a later publication.  \\subsubsection{Is the \\nvciv\\ Line Ratio a Reliable N/C Indicator?}  Hamann \\et (2002, hereafter H02) have used state-of-the-art photoionization calculations to investigate several line ratios in order to select those that are robust abundance indicators. The calculations span a vast range of densities [$7\\leq$ log $n_{\\rm H} (cm^{-3})\\leq 14$] that completely cover the typical range attributed to the broad emission line gas in AGNs [log $n_{\\rm H} (cm^{-3}) \\approx10$]. Since there is no evidence for densities as large as $10^{13}$--$10^{14} \\ cm^{-3}$ in NLS1s, we have no reason to suspect that they lie outside the range covered by the H02 calculations. According to the calculations, the \\nvciv\\ line ratio is a reliable N/C indicator over the range of interesting physical conditions expected in the BLR. H02 have also considered several of the weak intercombination lines, such as \\niiioiii\\ and \\nivciv\\ , previously discussed by Shields (1976), and concluded that the first is more reliable than the second, as it is less sensitive to the model assumptions. Results for all relevant line ratios are shown in their Figure~4 and a specific application to the locally optimally emitting clouds (LOCs) model is shown in H02 Figure~5.  The only other line ratio available to us, except for the \\nvciv\\ ratio discussed above, is \\nivciv. We have therefore investigated the N/C obtained from the two line pairs both observationally and theoretically. Several BLAGNs of our sample had published flux values (or upper flux limits) of the weak line \\ion{N}{4}]. We added to this subsample measurements of the \\ion{N}{4}] line in eight of our nine extreme NLS1s.  The N/C abundance obtained from the \\nivciv\\ line ratio was then calculated for the subsample, assuming that upper limits represent real ratios. We find that the two line ratios are well correlated ($r=0.8$ for 30 sources); i.e., if \\nvciv\\ is a good N/C indicator, so is \\nivciv. However, the derived N/C, {\\it assuming both ratios are reliable metallicity indicators}, is very different. The \\nvciv\\ ratio gives systematically larger N/C, sometimes by a factor as large as 3 or 4. We stress that the results are still tentative because of the large number of upper limits rather than real line ratios used in the analysis. Better data are required to confirm this correlation.  Regarding the suitability of \\nivciv\\ as an N/C indicator, we note that the calculations presented in Figure 4 of H02 show that the line ratio does not change by more than a factor of 2 over the range of conditions thought to be acceptable in AGN BLRs (ionization parameter of 0.03--0.3 for the H02 continuum and density below about 10$^{12}$ cm$^{-3}$). The H02 conclusion that the line ratio is not a robust N/C indicator is based on regions in parameter space that are different from the one specified above. Moreover, the particular example shown in Figure 5 of H02, applicable to the LOC model, clearly shows that under such conditions (that produce well most of the observed line ratios in AGNs; see Baldwin et al. 1995), the \\nivciv\\ line ratio is indeed a very good N/C indicator. A key issue is whether or not the physical conditions in the BLR of NLS1s are similar to those in broader line AGNs. The idea that density and optical depths may be different has been proposed in the past, and a better assessment of the \\nivciv\\ suitability in this case must await a more detailed theoretical investigation of such sources. At present we do not have a large enough sample and good enough calculations to test the suggestion that the results shown in Figure~2 are due to \\nvciv\\ being an inadequate N/C indicator.  \\subsubsection{High-Metallicity NLS1s? \\label{NLS1Z}}  The second scenario is based on the assumption that the \\nvciv\\ is a reliable N/C indicator. This led HF93 to suggest a strong $Z-L$ relationship in AGNs. However, our new measurements clearly show that NLS1s do not follow this $Z-L$ relationship. NLS1 metallicities, as indicated by the \\nvciv\\ ratio, are similar to those of the most luminous high-z quasars in our sample and are higher, by almost an order of magnitude, than those of BLAGNs with similar luminosities. In the \\nvheii\\ case, NLS1s show only slightly higher metallicities for a given luminosity compared with BLAGNs. This effect may be attributed to the more complex dependence of \\nvheii\\ on other physical parameters, such as the spectral energy distribution. Accepting this scenario, the HF93 $Z-L$ relationship cannot be a simple two-parameter dependence for all AGNs.  We also attempted to find a link between BLAGNs and NLS1s in order to see whether the $Z-L$ relationship is a smooth function of FWHM(H$\\beta$). We found no correlation between FWHM(H$\\beta$) and metallicity, luminosity, or a combination of the two. Since we have FWHM(H$\\beta$) values for about half of our sample, we checked whether FWHM(\\ion{C}{4}) could be more suitable, since this emission line is more dominant in our case. Again we found no correlation with the other parameters, nor any FWHM(\\ion{C}{4})--FWHM(H$\\beta$) correlation, for objects having both lines measured. In fact, we find that only AGNs that have FWHM(H$\\beta$)\\ltsim\\,1500 \\kms\\, follow a significantly different $Z-L$ relationship.  If high metallicity is indeed another NLS1 extreme property, the question whether this is related to some fundamental NLS1 physical parameter remains unanswered. The introduction of NLS1s to the \\nvciv\\ $Z-L$ diagram completely changes the correlation at low $L$ and implies that there is an additional dimension to this dependence that allows high metallicities at low luminosities. This hidden variable in the $Z-L$ relationship may be related to fundamental physical properties, such as the accretion rate or the age of the central BH. The key to answering this question possibly lies in the claim that NLS1s have low BH masses for their luminosities (equivalent to larger $L/L_{\\rm Edd}$ or higher accretion rate) and therefore follow a different mass-luminosity relationship than BLAGNs \\cite{pet00}. Boroson \\& Green (1992) pointed out that $L/L_{\\rm Edd}$ might be the underlying fundamental physical property of their primary eigenvector. Since NLS1s lie at one extreme end of this eigenvector, a natural hypothesis is that this property is also driving their unusually high metallicities. In order to test this hypothesis, one needs accurate BH mass determinations for our sample. Unfortunately, those are available for only 34 objects \\cite{kas00}, including only three extreme NLS1s. Moreover, measuring BH masses for high-z quasars is not a straightforward task since reliable determinations rely on reverberation mapping studies that would require at least a decade-long monitoring campaign due to cosmological time dilation. In addition, measurements of [\\ion{O}{3}], optical \\ion{Fe}{2}, and H$\\beta$ lines in high-z quasars are scarce (e.g. McIntosh \\et 1999).  \\subsection{The $Z-L$ Relationship in BLAGNs \\label{ZL}}  Finally, we remain within the framework of the second scenario, where the HF93 line ratios are considered reliable abundance indicators and NLS1s show systematically higher metallicities than BLAGNs with comparable luminosities.  We therefore removed all NLS1s from the sample and checked the $Z-L$ relationship for BLAGNs [we caution the reader that the BLAGN group includes objects with unknown FWHM(H$\\beta$); see \\S~\\ref{sample_prop}]. Our results confirm the observational evidence for a correlation between metallicity and luminosity (HF93) for a large (statistically incomplete) sample of BLAGNs. This correlation may be written in the form $Z \\propto L^{\\alpha}$, where $Z$ and $L$ are the BLR metallicity and BLAGN luminosity, respectively. Since both \\nvciv\\, and \\nvheii\\, are approximately proportional to $Z$ (HF93), the index is assigned with the mean value obtained for our sample, which is $\\alpha \\sim0.2$ (Tables 1 and 2). Combining this result with the empirical AGN BH mass-luminosity relationship \\cite{kas00}, $M_{\\rm BH} \\propto L^{0.5}$, we find $Z \\propto {M_{\\rm BH}}^{0.4}$. Alternatively, one may argue that the \\cite{kas00} result regarding the $L$ versus $M$ relation is biased because of the small size of the sample, and a more realistic form is perhaps $M_{\\rm BH} \\propto L$. In this case, $Z \\propto {M_{\\rm BH}}^{0.2}$. These results differ from those that were reached by supermassive BH growth considerations where $\\alpha$ was assumed to take a value of $\\approx\\,0.5$ \\cite{wan01}. We note, again, that the above obtained $\\alpha$ depends on the sample selection and is, therefore, rather uncertain."
    },
    "0201/astro-ph0201086_arXiv.txt": {
        "abstract": "We present dynamical models of the nearby compact elliptical galaxy M32, using high quality kinematical measurements, obtained with the integral-field spectrograph {\\tt SAURON} mounted on the William Herschel Telescope on La Palma. We also include {\\tt STIS} data obtained by Joseph et al. We find a best-fit black hole mass of $M_\\bullet = (2.5 \\pm 0.5) \\times 10^6 M_\\odot$ and a stellar $I$-band mass-to-light ratio of $(1.85 \\pm 0.15)\\,M_\\odot/L_\\odot$. For the first time, we are also able to constrain the inclination along which M32 is observed to $70^\\circ \\pm 5^\\circ$. Assuming that M32 is indeed axisymmetric, the averaged observed flattening of 0.73 then corresponds to an intrinsic flattening of $0.68 \\pm 0.03$.  These tight constraints are mainly caused by the use of integral-field data. We show this quantitatively by comparing with models that are constrained by multiple slits only. We show the phase-space distribution and intrinsic velocity structure of the best-fit model and investigate the effect of regularisation on the orbit distribution. ",
        "introduction": "M32 is a high-surface brightness, compact E3 companion of the Andromeda galaxy.  Ground-based kinematic measurements of this galaxy (Tonry 1984, 1987) already showed a steep gradient in the central velocity profile and a central dispersion peak, suggesting the presence of a central compact object, presumably a supermassive black hole. Since then, (spectroscopic) observations with increasing spatial resolution, both ground-based (Dressler \\& Richstone 1988; van der Marel et al. 1994a; Bender, Kormendy \\& Dehnen 1996) and with HST (Lauer et al. 1992; van der Marel, de Zeeuw \\& Rix 1997; Joseph et al.\\ 2001), strengthened the case for this black hole.  Simultaneously, dynamical models of ever improving quality were built to match these data sets (Richstone, Bower \\& Dressler 1990; van der Marel et al. 1994b; Qian et al.\\ 1995; Dehnen 1995). For most galaxies, the value of the best-fit black hole mass that is found depends on the assumptions that are made about the distribution function (DF) of the galaxy. Two-integral models, which assume that the DF depends only on the two classical integrals of motion (the energy $E$ and the vertical component of the angular momentum $L_z$), generally tend to overpredict the central black hole mass, since they cannot provide sufficient radial motion (Magorrian et al.\\ 1998; Gebhardt et al.\\ 2000; Bower et al. 2001).  The current state-of-the art models allow for the maximal degree of anisotropy by assuming that the distribution function depends on three integrals of motion (van der Marel et al.\\ 1998, hereafter vdM98; Cretton et al.\\ 1999, hereafter C99; Gebhardt et al. 2000, 2001). The internal kinematical structure of the best-fitting three-integral model of M32 closely resembles that of an $f(E,L_z)$ model. This explains why the central black hole mass that was found in M32 by the early models does not differ much from the current value of $(3.4 \\pm 0.7)\\times 10^6 M_\\odot$ (vdM98).  This means that the mass of the central black hole and intrinsic kinematical structure of M32 are constrained rather well. The main uncertainty that remains is the inclination along which M32 is observed, or, equivalently, the intrinsic flattening: edge-on models to a data-set consisting of {\\tt FOS} kinematics and ground-based observations along four position angles fit equally well as models that are projected over $55^\\circ$ (vdM98). When we combine this with the observed flattening of M32 (which is almost constant and equal to $0.73$ inside the central ten arcseconds, vdM98), we see that the intrinsic flattening of M32 is not very tightly constrained and can have any value between $0.55$ and $0.73$.  Recent developments in instrument design offer a way out: integral-field spectrographs capture the full two-dimensional behaviour of objects and are therefore expected to better constrain parameters such as the inclination. This tighter constraint on the intrinsic parameters can be explained from the behaviour of two-integral models, even though these generally provide less accurate fits to observational data. The part of a two-integral distribution function that is even in the velocities is only determined once the meridional plane density distribution is known (Qian et al.\\ 1995). The full intrinsic velocity map is needed to constrain the odd part of the DF, which means we need to measure the line-of-sight velocity $v_{\\rm los}$. Whereas these two lowest order velocity moments are sufficient to determine the {\\it stellar} DF of a galaxy, the dark matter distribution is only constrained once the velocity anisotropy is known (Dejonghe 1987; Gerhard 1991, 1993). This quantity can be measured through the higher order velocity moments. It is therefore necessary to determine the full two-dimensional kinematic behaviour of a galaxy. Capturing this information with only a few slit positions, as is generally attempted, is sometimes possible, but not always, which implies that the distribution function remains unconstrained.  These effects are even more important for {\\it three}-integral distribution functions, which, in most cases, provide better fits to the observations and have more freedom. Therefore, when galaxy models are constrained by kinematic observations along a modest number of slits, the intrinsic structure may remain largely unconstrained. In this paper, we revisit M32 using two-dimensional kinematical maps obtained with the panoramic integral-field spectrograph {\\tt SAURON} (Bacon et al.\\ 2001, hereafter Paper I; de Zeeuw et al.\\ 2002, hereafter Paper II) and high-resolution spectra obtained with {\\tt STIS} on board HST (Joseph et al.\\ 2001). We show that the use of integral-field data places very tight constraints on the central black hole mass and mass-to-light ratio of M32, but also, for the first time, on the inclination, or equivalently, the intrinsic flattening.  The paper is organized as follows: in Section \\ref{section2}, we give a brief summary of the data that are used to constrain the dynamical models, which are described in Section \\ref{section3}. The results are presented in Section \\ref{section4} and we show in Section \\ref{section5} that the use of integral-field data is crucial for most of the results that we obtained. The properties of the best-fitting model are described in Section \\ref{section6}, and we conclude with a discussion in Section \\ref{section7}.  Throughout the paper we assume a distance of $D=0.7$ Mpc (Welch et al.\\ 1986). This choice does not influence our conclusions, but sets the scale of the models in physical units. Specifically, lengths and masses scale as $D$, while mass-to-light ratios scale as $D^{-1}$. ",
        "conclusions": "\\label{section7}  We have presented dynamical models of the nearby compact elliptical M32, using data from the integral-field spectrograph {\\tt SAURON} and from {\\tt STIS} on board HST.  We have shown that our modeling software is able to deal with two-dimensional kinematical information and that the integral-field data tightens the constraint on all intrinsic model parameters considerably.  The axisymmetric three-integral model that best fits the data has a black hole mass of $(2.5 \\pm 0.5) \\times 10^6 M_\\odot$ and a stellar $I$-band mass-to-light ratio of $(1.85 \\pm 0.15)\\, M_\\odot/L_{\\odot}$. These values confirm the best-fit parameters that were obtained by previous authors, although our modeling procedure is different: we use a fully independent kinematical data-set and a different parameterisation for the mass density. Despite these differences in the methods, the best-fitting parameters that we find match the ones that were found by e.g. vdM98. This means that both approaches are reliable and that the results are robust.  For the first time, we are able to determine the inclination along which M32 is observed with great accuracy: the best-fit value is $70^\\circ \\pm 5^\\circ$.  If M32 is indeed axisymmetric, the averaged observed flattening of 0.73 then corresponds to an intrinsic flattening of $0.68 \\pm 0.03$. We have shown in Section \\ref{section5} that this tight constraint is mainly caused by the use of integral-field data. This implies that integral-field data will provide us with more insight into the internal structure and kinematics of such objects.  Although M32 is consistent with axisymmetry, it may be intrinsically triaxial, but seen along one of its principal planes. Allowing for an intrinsically triaxial object also would enable us to study the effects of uncertainties in the deprojection on the best-fitting parameters more closely. An extension of our method to triaxiality is in progress (Verolme et al.\\ 2002, in prep.), which will allow us to verify the assumptions made here. Furthermore, the tests that were carried out in this paper and the agreement with the results of other authors indicate that, given the assumptions, the results are robust.  We have obtained observations with {\\tt SAURON} of a representative sample of nearby ellipticals, lenticulars and spiral bulges, for most of which high-resolution {\\tt STIS} data is, or will become, available. We are in the process of applying the axisymmetric version of Schwarzschild's method on {\\tt SAURON} observations of a few of the sample galaxies that appear consistent with axisymmetry (e.g. NGC821, NGC3377, NGC2974).  The results of these dynamical models will provide us with the intrinsic parameters of a considerable number of objects, giving us unique insight into the formation and evolution of early-type galaxies.  \\vspace{0.5cm}  \\noindent{\\bf Acknowledgements}\\\\ We thank Karl Gebhardt, Richard McDermid and Glenn van de Ven for useful discussions, and Eric Emsellem, Harald Kuntschner and Reynier Peletier for a critical reading of the manuscript. YC acknowledges support through a European Community Marie Curie Fellowship. MB acknowledges support from NASA through Hubble Fellowship grant HST-HF-01136.01 awarded 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."
    },
    "0201/astro-ph0201279_arXiv.txt": {
        "abstract": "We present evidence for a correlation between expansion velocities of the ejecta of Type II plateau supernovae and their bolometric luminosities during the plateau phase. This correlation permits one to standardize the candles and decrease the scatter in the Hubble diagram from $\\sim$1 mag to a level of 0.4 and 0.3 mag in the $V$ and $I$ bands, respectively. When we restrict the sample to the eight objects which are well in the Hubble flow ($cz$ $>$ 3,000 km s$^{-1}$) the scatter drops even further to only 0.2 mag (or 9\\% in distance), which is comparable to the precision yielded by Type Ia supernovae and far better than the ``expanding photosphere method'' applied to Type II supernovae. Using SN~1987A to calibrate the Hubble diagrams we get $H_0$=55$\\pm$12. ",
        "introduction": "\\noindent Distances to cosmological objects are the path to get the expansion rate and the age of the universe. From observations of high-$z$ objects it is possible to measure how the expansion rate changes with time, and derive fundamental parameters like the geometry, deceleration, and energy content of the universe. This experiment has been recently done by two groups of astronomers using Type Ia supernovae (SNe) \\citep{riess98,perlmutter99}. Their observations revealed the surprising result that the universe is presently accelerating due to a non-zero cosmological constant, a form of dark energy that permeates space and dominates the total energy content of the universe, a result largely unanticipated by modern physics. If confirmed, the result from SNe~Ia would be a revolution in astrophysics. Before we can fully believe this result we need independent checks. Although Type II SNe are not as bright as the Ia's, they are the most common type of supernova and they offer the potential to be used as distance indicators using models of their atmospheres, a technique known as the ``expanding photosphere method'' (EPM) \\citep{schmidt94,hamuy01}. The scatter in the Hubble diagram shows that the precision in an EPM distance is $\\sim$20\\% \\citep{hamuy01}, which proves significantly larger than the 7\\% precision yielded by SNe~Ia \\citep{hamuy96,phillips99}, thus hampering the use of SNe~II for determination of cosmological parameters.  In principle, the apparent magnitudes of stellar objects can be used to derive distances, as long as a class of objects with known luminosities can be identified. Although Type II SNe display a wide range in luminosities at all epochs -- making it hard to use them as standard candles -- in this letter we show that the envelope expansion velocities of the plateau subclass \\citep{barbon79} are highly correlated with their luminosities during the plateau phase. This finding allows us to standardize the candles and use them to derive distances with precisions comparable to that delivered by SNe~Ia. ",
        "conclusions": ""
    },
    "0201/astro-ph0201109_arXiv.txt": {
        "abstract": "Detailed investigations of extensive air showers have been performed with the data measured by the KASCADE experiment. The results allow to evaluate hadronic interaction models, used in simulations to interpret air shower data. The all-particle spectrum of cosmic rays and their mass composition, as well as individual spectra for groups of elements have been reconstructed. The results suggest, the {\\sl knee} in the all--particle cosmic--ray energy spectrum is caused by a rigidity--dependent cut--off of individual element groups. \\vspace{1pc} ",
        "introduction": "The earth's atmosphere is continuously bombarded by highly relativistic ionized particles, first discovered and named \"cosmic rays\" by V.~Hess in 1912. Present--day experiments show the cosmic--ray energy spectrum extending up to more than $10^{20}$ eV. The flux spectrum follows a power law $dN/dE \\propto E^{-\\gamma}$ over many decades in energy. The most prominent feature is the {\\sl knee} in the spectrum around 3~PeV where the spectrum steepens from $\\gamma\\approx 2.7$ to $\\gamma\\approx 3.1$. The origin of cosmic rays is still under debate. Strong, relativistic shock fronts expanding from supernova explosions are favoured by popular models for the acceleration of ionized particles. Such models explain the particle acceleration up to energies of about $Z\\cdot 10^{15}$ eV, with the nuclear charge $Z$ of the particle. This upper limit coincides for primary protons with the mentioned steepening of the spectrum, and the origin of the {\\sl knee} is related to the upper limit of acceleration in several models.  Since the charged particles are deflected in the interstellar magnetic fields, the only hint of their sources are their energy spectrum and the mass composition, or more preferable, the energy spectra of individual elements. Cosmic rays at energies below 1~PeV have been directly observed by balloon--borne instruments at the top of the atmosphere or in outer space. At higher energies, the steep falling flux spectrum requires large detection areas or long observation periods, presently only possible in ground--based installations. These detector systems measure the secondary particles produced by cosmic rays in the atmosphere, the extensive air showers (EAS).  To investigate the cosmic rays from several $10^{13}$ eV up to $10^{17}$ eV the air shower experiment KASCADE (\"Karlsruhe Shower Core and Array DEtector\") \\cite{kascade} has been built on--site at the Forschungszentrum Karlsruhe in Germany. The experiment detects the three main components of EAS simultaneously. A $200\\times 200$ m$^2$ scintillator array measures the electromagnetic and muonic components. The 320 m$^2$ central detector system combines a large hadron calorimeter \\cite{kalo} with several muon detection systems \\cite{mwpc}. In addition, high energetic muons are measured by an underground muon tracking detector \\cite{muontunnel}. ",
        "conclusions": "The EAS observables obtained by KASCADE are sensitive to hadronic interaction models used in simulations to interpret EAS data. At present, the combination CORSIKA/QGSJET best describes the measurements. The systematic dependence of the mean logarithmic mass on different observables and models has been investigated. The systematic error for the model QGSJET is in the order of $\\Delta\\langle\\ln A\\rangle\\approx0.8$. The reconstruction of individual energy spectra for groups of elements indicate a rigidity dependent cut--off, which explains the {\\sl knee} in the all particle cosmic--ray flux spectrum.  { The KASCADE experiment is supported by the Ministry for Research of the German government and embedded in collaborative WTZ projects between Germany and Romania (RUM 97/014), Poland (POL 99/005), and Armenia (ARM 98/002). The Polish group acknowledges the support by KBN grant no. 5\\,P03B\\,133\\,20. }"
    },
    "0201/astro-ph0201423_arXiv.txt": {
        "abstract": "We present results from the combination of two \\chandra ~ pointings of the central region of the cluster of galaxies A3667. From the data analysis of the first pointing Vikhlinin et al. reported the discovery of a prominent cold front which is interpreted as the boundary of a cool gas cloud moving through the hotter ambient gas. Vikhlinin et al. discussed the role of the magnetic fields in maintaining the apparent dynamical stability of the cold front over a wide sector at the forward edge of the moving cloud and suppressing transport processes across the front. In this Letter, we identify two new features in the X-ray image of A3667: i) a 300 kpc arc-like filamentary X-ray excess extending from the cold gas cloud border into the hotter ambient gas; ii) a similar arc-like filamentary X-ray depression that develops inside the gas cloud.  Both features are located beyond the sector identified by the cold front and are oriented in a direction perpendicular to the direction of motion. The temperature map suggests that the temperature of the filamentary excess is consistent with that inside the gas cloud while the temperature of the depression is consistent with that of the ambient gas.  We suggest that the observed features represent the first evidence for the development of a large scale hydrodynamic instability in the cluster atmosphere resulting from a major merger. This result confirms previous claims for the presence of a moving  cold gas cloud into the hotter ambient gas. Moreover it shows that, although the gas mixing is suppressed at the leading edge of the subcluster due to its magnetic structure, strong turbulent mixing occurs at larger angles to the direction of motion. We show that this mixing process may favor the deposition of a nonnegligible quantity of thermal energy right in the cluster center,  affecting the development of the central cooling flow. ",
        "introduction": "The central region of A3667, a nearby, hot merging cluster (Markevitch, Sarazin, \\& Vikhlinin 1999), was observed for the first time by \\chandra ~ in Sept 1999. The analysis of this first observation by Vikhlinin, Markevitch \\& Murray (2001a,b) reveals the presence of a prominent 500~kpc-long density discontinuity (``cold front'') in the cluster atmosphere. Vikhlinin et al.\\ show that: i) the density discontinuity is the contact surface between the moving cloud of 4~keV gas and the hotter ambient gas; ii) the speed of the moving cloud is slightly supersonic.  An interesting aspect raised by Vikhlinin et al. is that the cold front must develop hydrodynamical instabilities. These instabilities would destroy the front on very short time scales. However, the \\chandra ~ image shows that the front is stable in a wide, $\\varphi =\\pm 30^\\circ$ sector (where $\\varphi$ is the angle with respect to the direction of motion). Moreover, the front width is smaller than the Coulomb mean free path, which indicates that the transport and mixing processes are suppressed (see also Ettori \\& Fabian 2000). To reconcile this observational evidence with the expected hydrodynamical instability of the cold front, Vikhlinin et al. (2001b) put forward the following scenario. As the cloud moves through the ambient gas, the magnetic field lines, which are frozen in the intracluster medium, are stretched by the tangential plasma motions. This results in the formation of a layer with the ordered magnetic field parallel to the front surface. The magnetic field intensity in such a layer is sufficient to suppress the hydrodynamical instabilities within the $30^\\circ$ sector, and to prevent gas mixing and transport processes.    \\begin{inlinefigure} \\centerline{\\includegraphics[width=0.95\\linewidth]{f1.eps}} \\caption{Count rate image in the $0.7-4$~keV band binned by $4^\\second$. The arrows indicate the prominent X-ray features: the cold front discussed by Vikhlinin et al. (2001a,b), a filamentary arc-like surface brightness excess extending toward the east, and a filamentary arc-like surface brightness depression extending toward the west.} \\label{fig:image} \\end{inlinefigure}   In this letter we present the results from the joint analysis of two \\chandra ~ observations of A3667. We identify two filamentary structures forming at the border of the merging subclump: one extends toward the outskirts, the other toward the cluster center. Both structures lie well beyond the  hydrodynamically stable region identified by Vikhlinin et al. (2001a). We speculate that these filaments correspond to the well-developed K-H instability, and discuss their implications for the proposed dynamical model of A3667.   The physical size of the structures under discussion is computed assuming $H_0=50$~km~s$^{-1}$~kpc$^{-1}$ ($1$~arcsec$\\approx 1.46$~kpc). Unless specified differently, all the errors are at $90\\%$ confidence level for one interesting parameter.    \\begin{inlinefigure} \\centerline{\\includegraphics[width=0.95\\linewidth]{f2.eps}} \\caption{Adaptively smoothed \\chandra ~ image. The four lines identify the External Sector (ES) and the Internal Sector (IS) used to derive the surface brightness and temperature profiles reported in Fig.~\\ref{fig:p_dep}. The position angles are from -15\\degd ~ to 15\\degd ~ and from -15\\degd ~ to 40\\degd ~ for IS and ES, respectively (angles are measured from North through East). Regions 4 and 2 identify the X-ray excess and depression, respectively. Regions 1 and 3 are two circular $r=200$~kpc regions centered at the center of curvature of each filamentary structure. The actual temperatures of the regions, indicated by cardinal numbers 1 to 4 are, reported in Fig.~\\ref{fig:t_reg}.} \\label{fig:panda} \\end{inlinefigure} ",
        "conclusions": "We presented the combination of two \\chandra ~ observations of the central region of A3667.  We showed that the cluster hosts two arc-like adjacent filamentary structures: one, extending from the colder subcluster toward the cluster outskirts, appears as a dense structure embedded in the less dense cluster atmosphere, and the other, extending inside the subcluster, appears as a rarefied structure embedded in the denser cluster core.  We suggest that the observed features represent the first evidence for the development of a large scale hydrodynamic instability in the cluster atmosphere.  This interpretation appears to be consistent with the previous cluster dynamic interpretation proposed by Vikhlinin et al. (2001a,b).  The discovery of this instability represents an important step toward a full understanding of the physics of mergers. In particular it shows that, although the cold front prevents the gases of the merging objects to mix along the direction of motion, strong turbulent mixing processes, on scales comparable to the size of the merging subclump, may occur at large angles to the direction of motion. This may favor the deposition of a non-negligible quantity of thermal energy right in the cluster center  with important consequences for the development and/or evolution  of a central  cooling flow."
    },
    "0201/cond-mat0201002_arXiv.txt": {
        "abstract": "\\ \\ \\ Complex adaptive systems have been the subject of much recent attention. It is by now well-established that members (`agents') tend to self-segregate into opposing groups characterized by extreme behavior. However, while different social and biological systems manifest different payoffs, the study of such adaptive systems has mostly been restricted to simple situations in which the prize-to-fine ratio, $R$, equals unity. In this Letter we explore the dynamics of evolving populations with various different values of the ratio $R$, and demonstrate that extreme behavior is in fact {\\it not} a generic feature of adaptive systems. In particular, we show that ``confusion'' and ``indecisiveness'' take over in times of depression, in which case cautious agents perform better than extreme ones. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0201/astro-ph0201490_arXiv.txt": {
        "abstract": "We present the study of ten random realizations of a density field characterized by a cosmological power spectrum $P(k)$ at redshift $z=50$. The reliability of such initial conditions for \\nbody\\ simulations are tested with respect to their correlation properties. The power spectrum $P(k)$, and the mass variance $\\sigma_M(r)$ do not show detectable deviations from the desired behavior in the intermediate range of scales between the mean interparticle distance and the simulation volume. The estimator for $\\xi(r)$ is too noisy to detect any reliable signal at the initial redshift $z=50$. The particle distributions are then evolved forward until $z=0$. This allows us to explore the cosmic variance stemming from the random nature of the initial conditions. With cosmic variance we mean the fact that a simulation represents a single realization of the stochastic initial conditions whereas the real Universe contains many realizations of regions of the size of the box; this problem affects most importantly the scales at about the fundamental mode. We study morphological descriptors of the matter distribution such as the genus, as well as the internal properties of the largest object(s) forming in the box. We find that the scatter is at least comparable to the scatter in the fundamental mode. ",
        "introduction": "\\label{Intro}  Our present understanding of the formation and properties of the cosmological large-scale structure relies to a large extent on \\nbody\\ simulations: given the difficulty in addressing theoretically the highly nonlinear regime of the growth of density inhomogeneities by the gravitational instability, simulations have proven a valuable tool to get insight into the (non-linear) structure formation scenarios. Therefore, it is of considerable importance to confirm the reliability of such simulations.  It has been claimed recently (Baertschiger~\\& Sylos Labini 2002) that there are major problems with generating initial conditions (ICs) for \\nbody\\ simulations. We can identify several reasons why the ICs {\\em may} introduce uncertainties in the subsequent evolution. First, there is the problem of finite-mass resolution or {\\em discreteness}: the initial continuum density field is modeled by the distribution of a {\\em finite} number of point particles $N$, therefore only a finite number of Fourier modes of the density field can be reproduced reliably. The maximum wavenumber ({\\it Nyquist wavenumber}) is given by $k_{\\rm max} = \\pi/\\Delta x$, where $\\Delta x$ is the mean interparticle separation.  The modes $k \\gtrsim k_{max}$ have spurious values related to the point-particle distribution and may lead to artificial effects in the posterior dynamical evolution.  The finite-mass resolution is expected to be irrelevant if the nonlinear mode-mode coupling to the modes $k \\gtrsim k_{max}$ has only a small influence on the dynamics.  The second problem with the ICs is the finite size of the simulation box with side length $B$, which implies that the values of the Fourier modes of the density field with wavenumber smaller than the {\\it fundamental wavenumber}, $k < k_{\\rm min} = 2\\pi/B$, are artificially set to zero. This leads to two possible difficulties: first, the absence of mode-mode coupling to those large-scale modes, and second the so-called {\\em cosmic variance}, meaning that the simulation box represents only one (finite-sized) realization of the stochastic initial density field, whereas the true Universe contains many realizations of regions of the size of the box. Therefore, the morphological properties of the matter distribution in a certain volume, as measured by e.g. the genus statistics, will presumably show some intrinsic scatter when placing the volume at different locations in the real Universe. And this will also happen with the (internal) properties of any given class of objects. This is one of the main aspects of the current study and what we refer to as cosmic variance (in \\nbody\\ simulations) throughout the paper even though one might argue that this is not the \"real\" cosmic variance but rather an artificially introduced \\textit{sampling variance}. However, we are actually interested in exactly that (sampling) effect which can easily be tested by simply running the same cosmological simulation but using different random realizations of the initial density field.  In this work we study systematically the reliability of the initial density field used as an input to the \\nbody\\ simulations as well as the effect of their random nature onto the internal properties of clusters. In \\Sec{nbody} we briefly review the most commonly way used to generate ICs and the code used to evolve the particles into the non-linear regime. \\Sec{DManalysis} focuses on some of the statistical characteristics of the Dark Matter field: we consider the 2-point correlation function, the power spectrum, the mass variance in spheres, and the Minkowski functionals, the latter being sensitive to correlations of higher order. Finally, in \\Sec{HaloAnalysis} we investigate dark matter clusters identified within the simulations and quantify the effect of the cosmic variance on their internal properties. ",
        "conclusions": "We presented the study of ten random realizations of a density field characterized by a cosmological power spectrum $P(k)$ at redshift $z=50$. These initial conditions for \\nbody\\ simulations were tested with respect to their correlation properties. Recent claims by Baertschiger~\\& Sylos Labini (2002) throw doubts on the ability of the commonly used method for generating the initial density field using particles (i.e. glass or grid preinitial distribution + the Zeldovich approximation, Eftstahiou~\\ea 1985) to clearly reproduce the analytical input correlations. The power spectrum $P(k)$ and the mass variance $\\sigma_M(r)$ do not deviate from the expected behavior (including expected departures from the desired \\LCDM\\ behavior due to finite mass and finite size effects). The estimated 2-point correlation $\\xi(r)$ is too noisy to be used as a reliable credibility check; one cannot claim either that it reproduces the desired \\LCDM\\ behavior or that it exhibits systematic deviations thereof.  These initial conditions were then evolved forward in time until redshift $z=0$ using the publicly available adaptive mesh refinement code \\mlapm\\ (Knebe, Green~\\& Binney 2001). This allowed us to explore the cosmic variance stemming from the random nature of the initial conditions, i.e., the scatter between different realizations of statistically identical initial conditions. We addressed the morphological properties of the matter distribution with the four Minkowski functionals as functions of a density threshold. The scatter grows in time, the one exhibiting a larger dispersion being the genus, of the order of $10\\%$ at $z=0$.  We also investigated the internal properties of DM halos, which have already been shown by other groups to be profoundly influenced by the surrounding large-scale structure, which in turn is sensitive to $k$-modes $\\approx$ fundamental mode (Colberg~\\ea 1999). We find that the scatter in the properties of the most massive object(s) forming in the box is $\\sim 20\\%$, and as high as $\\sim 50\\%$ for some properties such as the mass or the spin parameter.    An interesting question is whether this scatter is induced mainly by the cosmic variance of the amplitude at scales around the fundamental mode, or by the cosmic variance of the random phases. There is certainly a propagation of the error in the initial large-scale amplitude by power transfer towards smaller scales. In fact, when comparing our data to the (non-)linear fit of Peacock~\\& Dodds (1996) for the power spectrum and to the prediction by Press~\\& Schechter (1974) for the mass function, we find good agreement. The derivation of both results is based on the hypothesis of small influence from coupling of modes at some $k$ to modes with larger $k$; our results support this assumption, as far as the statistical estimators we probed are concerned. It would now be interesting to investigate in detail the actual influence of the large waves onto the small scale structure. This would also shed some light on the credibility of running small simulation boxes to very low redshifts as already done by several groups (e.g. Dave~\\ea 2001, Avila-Reese~\\ea 2001, Gnedin 2000, Colin, Avila-Reese~\\& Valenzuela 2000), but we leave this to a future study."
    },
    "0201/astro-ph0201173_arXiv.txt": {
        "abstract": "The first galaxies formed at high redshifts, and were likely substantially less massive than typical galaxies in the local universe.  We argue that (1) the reionization of a clumpy intergalactic medium (IGM) by redshift $z\\approx 6$, (2) its enrichment by metals by $z\\approx 3$ without disturbing the Ly$\\alpha$ forest, and (3) the presence of supermassive black holes powering the recently discovered bright quasars at $z\\sim 6$, strongly suggest that a population of low--mass galaxies exists beyond redshifts $z\\gsim 6$. Although the first stars could have been born in dark matter halos with virial temperatures as low as $T_{\\rm vir}\\approx 200$K, collapsing as early as $z\\sim 25$, the first galaxies likely appeared in significant numbers only in halos with $T_{\\rm vir}>10^4$K that collapsed later ($z\\sim 15$).  The gas in these more massive halos initially contracts isothermally to high densities by atomic Ly$\\alpha$ cooling. ${\\rm H_2}$ molecules can then form efficiently via non--equilibrium gas--phase chemistry, allowing the gas to cool further to $T\\sim 100$K, and fragment on stellar mass scales.  These halos can harbor the first generation of ``mini-galaxies'' that reionized the universe.  The continuum and line emission from these sources, as well as their Ly$\\alpha$ cooling radiation, can be detected in the future by {\\it NGST} and other instruments. ",
        "introduction": "Recent measurements of the cosmic microwave background (CMB) temperature anisotropies, determinations of the luminosity distance to distant type Ia Supernovae (SNe), and other observations have led to the emergence of a robust ``best--fit'' cosmological model with energy densities in cold dark matter (CDM) and ``dark energy'' of $(\\Omega_{\\rm M},\\Omega_{\\rm \\Lambda})\\approx (0.3,0.7)$.  The growth of density fluctuations, and their evolution into non--linear dark matter structures can be followed in detail from first principles by semi--analytic methods~\\cite{13,14} and N--body simulations~\\cite{15}. Structure formation is ``bottom--up'', with low--mass halos condensing first. Halos with the masses of globular clusters, $10^{5-6}\\msun$, are predicted to have condensed as early as $\\sim$1\\% of the current age of the universe, or redshift $z\\sim 25$.  It is natural to identify these condensations as the sites where the first astrophysical objects, such as stars, or quasars, were born. ",
        "conclusions": ""
    },
    "0201/astro-ph0201459_arXiv.txt": {
        "abstract": "Using the VLA\\footnote{The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.}, we have made spectral-line and continuum observations of the neutral hydrogen in the direction of the compact group of galaxies Stephan's Quintet.  The high-velocity clouds between 5600 and 6600 \\kms, the disk of the foreground galaxy, NGC 7320, at 800 \\kms, the extended continuum ridge near the center of the group, and 3 faint dwarf-like galaxies in the surrounding field were imaged with C, CS, and D arrays.  Four of the HI clouds previously detected are confirmed. The two largest HI features are coincident with and concentrated mainly along separate large tidal tails that extend eastward. The most diffuse of the four clouds is resolved into two clumps, one coincide with tidal features south of NGC 7318a and the other devoid of any detectable stellar or \\Ha\\ sources. The two compact clouds, along the same line of sight, have peak emission at luminous infrared and bright \\Ha\\ sources probably indicative of star-forming activity. The total amount of HI detected at high redshifts is $\\sim 10^{10}M_\\odot$. As in previous HI studies of the group, no detectable emission was measured at the positions of any high-redshift galaxies so that any HI still bound to their disks must be less than $2.4 \\times 10^{7}M_\\odot$. ",
        "introduction": "\\label{sec:intro}  Stephan's Quintet (Arp~319; ``SQ\" hereafter) was the first compact group to be discovered in the late 1800's, and it is probably the best known and most studied dense group at all wavelengths.  The original membership identified by \\citet{s77} included five, bright, relatively isolated galaxies (NGC 7317, 7318a\\&b, 7319, and 7320) in a tight configuration on the sky.  In addition to their close proximity, three of the spiral members (NGC 7318a\\&b, 7319) have peculiar or highly distorted optical images that include faint wisps, filaments, and tails.  Given its high surface brightness \\citep[22.3 mag arcsec$^{-2}$;][]{Hickson82} and surface density enhancement \\citep[593;][]{sul87}, it is no surprise to find the quintet included among the Hickson groups \\citep{Hickson82,Hickson93}. While the physical compactness of most HCG's have been challenged \\citep[see][and references therein]{Barns96}, the optical evidence of recent tidal interactions leaves little doubt that some of the galaxies in the quintet are physically close.  As is typical of Hickson groups \\citep{sul97}, the quintet contains a discordant-redshift member (NGC 7320).  A difference in redshift of nearly 6000 \\kms\\ between NGC 7320 and the high-redshift members was first reported by \\citet{burbidge61} who concluded that NGC 7320 is either a member of the dense group or a foreground galaxy. The discovery of the discordant redshift in SQ and similar groups such as Seyfert's Sextet and VV172 \\citep{VV59} sparked much debate concerning the relationship between the high- and low-redshift galaxies in compact configurations and non-cosmological redshifts.  Initially, 21-cm line studies of SQ were undertaken mainly with the hope that the radio observations would provide distances independent of redshift, thereby resolving both questions about the group's membership and the existence of non-cosmological redshifts. \\citet{Allen70} was the first to detect HI emission in the direction of the quintet.  This emission, detected at 800 \\kms, was associated with NGC 7320.  Using the integral properties derived from the HI profile and comparing them with similar properties of field galaxies, \\citet{Allen70} concluded that NGC 7320 is a dwarf galaxy at a distance consistent with its redshift.  Because Allen's study did not resolve the question of group membership, \\citet{Balkowski73} and \\citet{Shostak74} were motivated to search for HI emission from other quintet galaxies that possibly could.  Both detected HI emission near 6600 \\kms\\ which they attributed to NGC 7319.  After applying the same method used by  \\citet{Allen70}, they derived very different distances. \\citet{Balkowski73}  placed the galaxy at the same distance as NGC 7320 while  \\citet{Shostak74} found it more likely that NGC 7319 belongs at a distance commensurate with its redshift.  The difference in their distances can be traced to the calibration samples that each adopted.  The main weakness of this method used to derive distances was the underlying assumption that the galaxies in the quintet have normal HI properties for their morphological types and luminosity.  Using the Westerbork Synthesis Radio Telescope (WSRT), \\citet{Allen80} would test this assumption by providing new information on the distribution and kinematics of the HI gas in the quintet.  Their observations supported the assumption of normalcy in the location and motion of the low-redshift HI but raised serious doubt about the use of HI profiles as distance indicators when the data showed the gas at 6600 \\kms\\ in an extended cloud (2\\arcmper 5 $\\times$ 4\\arcm) not directly associated with any of the members.  The discovery of an isolated cloud in SQ was significant because it, alone, changed the focus of subsequent 21-cm line studies.  HI observations more sensitive to extended emission over larger regions of the quintet were undertaken in an effort to understand this anomalous feature in the context of the group's origin and evolution.  Sensitive HI searches that used the Arecibo telescope found the presence of low level (8 mJy) broadband emission between 5600 and 6600 \\kms\\ \\citep{s80} and two more features extended and offset from the quintet galaxies \\citep{Peterson80}.  With better angular resolution and sensitivity than the old observations \\citep{Allen80}, \\citet{Shostak84} used the WRST to locate and confirm the three cloud features at 5700, 6000 and 6600 \\kms\\ that had been previously detected.  Their synthesis observations showed the HI largely outside the optical boundaries of the galaxies in the quintet.  This present study of SQ is part of a broader investigation of the distribution and kinematics of HI in the least and most compact associations in the Hickson catalog \\citep{will99a,will99b,vm00a,vm00b,vm01,hu00} in order to analyze the type and effect of the interactions that are taking place. If physically dense, Hickson groups are natural sites for studying tidal interactions and its effects, galaxy formation and evolution via merging processes, and the dynamical evolution of galaxy groups. Spiral galaxies normally more extended in HI than in starlight are susceptible to tides and direct collisions and ought to be more sensitive to recent interactions.  High resolution imaging of the neutral hydrogen can be used to distinguish between the real and illusory compact groups and may provide a sensitive test of the dynamical evolution of the physically dense groups. We have analyzed the spatial distributions and kinematics of HI within a subset of 16 groups and proposed an evolutionary scenario where  the amount of detected HI would decrease further with evolution, either by continuous tidal stripping and heating or by shocks produced when galaxies penetrate the group's core and interact violently with the interstellar or intergalactic HI in the groups \\citep{vm01}. SQ, amongst the densest Hickson groups, is a vital constituent in this scenario, hence its detailed study  would be essential in characterizing the extreme stages in galaxy interactions and group evolution.  There is strong observational evidence that SQ is a physical system.  The filaments, tails, and peculiar images of the galaxies are suggestive of interactions.  The distribution of the HI outside the galaxies is consistent with a history of collisional encounters that may have stripped the spiral members of their HI disks \\citep{Allen80,Shostak84}. Why then observe SQ with the VLA?  What more had we hoped to learn from HI synthesis observations? The WRST \\citep{Shostak84} integrated map gives us a clear sense of how the gas between 5700 and 6600 \\kms\\ is distributed.  Given the higher angular and velocity resolution of the VLA, we saw an opportunity to improve the kinematical description of the HI gas already detected outside of the galaxies.  With its higher sensitivity to faint compact emission, the VLA could detect additional isolated HI clouds within the quintet and possibly resolve any weak HI emission associated with tidal features or still bound to the disks of the individual galaxies. ",
        "conclusions": "\\label{sec:discussion}  New synthesis observations have motivated us to re-assess previously published results and revise dated images that have characterized the HI morphology of SQ for nearly twenty years.  The new observations at higher angular and velocity resolutions than any published results to date provide far more details about the internal motions of the HI gas in SQ and allow direct comparisons between HI emission and other discrete and extended sources within the group over a range of wavelengths, i.e., X-ray, optical, \\Ha, continuum, and infrared.  The higher resolution of the VLA has provided us with additional clues about the short-term consequences of collisions and interactions as well as the dynamical history of the high-redshift members in the quintet.  New details on substructures within all HI features detected, including the foreground galaxy NGC~7320, and the HI cloud at 6600 \\kms, hold the valuable clues about the dynamical history of the group.   \\subsection{Gaseous Collisions in the Group}  The large velocity difference of NGC~7318b with respect to the rest of the group ($\\Delta V \\sim 700$ \\kms) has been taken as key evidence for its being an intruder passing through the group.  A collision at such a high velocity would result in wide-spread shocked and ionized regions, followed by rapid cooling due to the emission of radiation at UV- and X-ray wavelengths \\citep[e.g.][]{Harwit87}. The radio continuum ridge found east of NGC~7318b has been suggested to be the interface between the interloper and the group because the shocked, ionized gas is predicted to produce significant amounts of synchrotron emission \\citep{vdHulst81,Allen72}. The synchrotron life time for the radiating electrons is estimated to be rather short, a few times $10^7$ years, which is comparable to the crossing time for such a collision \\footnote{A classic example of enhanced radio synchrotron emission arising from a collision involving two galaxies is the ``Taffy'' pair \\citep[UGC~12914/5;][]{Condon93}.  The HI observations of this colliding system should also help in understanding the response of cold gas within the progenitor's disks (see below).}. Presence of extended, diffuse X-ray emission associated with the group \\citep{Sulentic95,Pietsch97} also raises the possibility that what is seen in SQ is the result of a collision between a galaxy and a hot group medium.  Given the $n_e^2$ dependence on emissivity, such a collision should be much less effective in producing radio continuum emission than in a collision involving two galaxy disks.  A logical place to look for such an example is the centers of nearby clusters such as Virgo and the Coma cluster, and NGC~4522 may be an example of such a collision between diffuse cluster/group medium and a galactic disk \\citep{Kenney99,Kenney00}. Though plausible, such objects are rare in clusters, and the tenuous intra-group medium in SQ would weaken the case considerably.  The excellent surface brightness sensitivity achieved by our new observations reveals several new details for the first time, including the faint structure in the extended radio ridge between NGC~7318b and NGC~7319. As seen in Figure~\\ref{fig:cont}, the continuum ridge approximately follows the optical tails. The new continuum image also raises minor issues for the hypothesis that the extended radio continuum ridge, also visible in X-ray and \\Ha, is the product of a collision involving NGC~7318b. Because NGC~7318b sits in the middle of the ridge its trajectory has to be mainly in the direction perpendicular to the extent of the ridge; however, the shape of the radio ridge is S-shaped and not the expected symmetric C-shape of a bow-shock.  A careful comparison of radio continuum and \\Ha\\ images by \\citet{Ohyama98} shown in Figure~\\ref{fig:ridge} reveals some important differences. The morphology of the \\Ha\\ emission is indeed C-shaped, more in line with the shock hypothesis. This apparent discrepancy may be resolved at least in part by the difference in the radiation time scales.  The cooling time for the \\Ha\\ emitting gas is much shorter than the synchotron lifetime for the radio emitting electrons, and \\Ha\\ is tracing only the most active, ongoing shocked regions. There is also some evidence that more than one source of \\Ha\\ emission may be present.  The continuum subtracted line image centered at 6738 \\AA\\ (including both \\Ha\\ and [N II] emission) by \\citet{Sulentic01} resembles our radio continuum image more closely, and the radio continuum may be tracing mainly the ionized gas at a the higher velocity (i.e. 6600 \\kms) as shown in Figure 2 of \\citet{Xu99}.  The striking morphological alignment between \\Ha\\ and the HI features near 6600 \\kms\\, shown on the right panel in Figure~\\ref{fig:ridge} raises an intriguing possibility that NGC~7318b may be colliding with part of the HI tidal tail of Arc-S that runs behind NGC~7320.  Implicit in this hypothesis is that this HI tail may be linked with the NW-HV feature, and the missing segment in HI may be fully ionized and visible only in \\Ha\\ and in radio continuum.  The recession velocity of the ionized gas in the southern parts of the \\Ha\\ filament measured by \\citet{Ohyama98} is around 6600 \\kms, consistent with this scenario.  Further, the broad linewidth ($\\Delta v \\sim 900$ \\kms) measured by \\citet{Ohyama98} is also consistent with such a collision scenario.  \\subsection{Star Formation Activities in SQ \\label{sec:SF} }  Examples of close correspondence between \\Ha\\ emission and HI are found throughout SQ (see \\S~\\ref{sec:HIdist}).  The two prominent star-forming complexes identified in mid-IR by \\citet{Xu99} are associated with two bright \\Ha\\ complexes, and they both coincide with the two highest HI column density peaks (Arc-N and NW-HV). The HI peaks in the NW-HV, NW-LV, and the SW features all have associated \\Ha\\ features, and the molecular gas complexes traced in CO by \\citet{Yun97} and by \\citet{Gao00} are also associated with \\Ha\\ features detected by \\citet{pl99}, \\citet{Moles98}, \\citet{Ohyama98}, \\citet{Sulentic01}, and others.  Frequent tidal encounters in high density regions are thought to be conducive to star-formation activity, and indeed evidence for induced star-formation activity is ubiquitous everywhere in SQ where cold gas is found.  The star-formation activity, however, is not as spectacular as one might expect as in isolated environments such as in interacting pairs. Star formation is known to have a non-linear dependence on gas density \\citep[e.g.][]{Kennicutt98}, and few high density regions are currently present in SQ and little cold gas is remaining in the individual galaxies.  These trends of low gas content and ubiquitous but suppressed star-formation activity are commonly found in our examination of a large sample of Hickson Compact Groups \\citep{vm01}.   \\subsection{Tidal Features and Interaction Scenarios}  The improved sensitivity and resolution of the new VLA data have significantly improved our understanding of the nature and distribution of the HI tidal features in SQ (see \\S~\\ref{sec:HIdist}).   The massive L-shaped HI arc located about $3'$ (75 kpc) east of SQ is shown to be possibly consisting of two separate tidal features (Arc-N and Arc-S in Fig.~\\ref{fig:arc}), in accordance with the two optical tails seen. Most of the gas in Arc-N is located near the end of the long sinuous tail emanating from NGC 7319. While Arc-N appears to originate from NGC~7319, the origin of the southern tail is uncertain.  Both of the NW clouds are located near the tips of 2 crossing optical features north of NGC 7318a\\&b. The two NW features are quite distinct in velocity, but they appear spatially coincident exactly within our resolution. Therefore the two may be physicaly related.  The SW feature at 5700 \\kms\\ is shown to be a distinct feature consisting of 2-3 clumps.  If only its velocity field is considered, it might be an independent dynamical entity, such as a dwarf galaxy (see \\S~\\ref{sec:SW}).  One of the clear conclusions that can be drawn from the new data is that the 5700 \\kms\\ feature and 6000 \\kms\\ feature are clearly not the tidally disrupted disk of NGC~7318b as proposed by \\citet{Moles98}. In fact, none of the HI features are clearly associated with any of the individual galaxies. It is conceivable that in the past the two features could have been part of a single structure, possibly related to NGC~7318b, and separated into two distinct features forming the components we see today.  Such a separation cannot be a direct outcome of a tidal disruption because {\\it everything} within the tidal radius, including gas and stars, should have responded identically to the tidal shear.  To the contrary, the core of the optical galaxy appears intact.  The suggestion that NGC~7318b may be passing through the group for the first time \\citep{Moles98,Sulentic01} is certainly unlikely, especially if the observed X-ray, \\Ha, and radio continuum ridge is the evidence for the ongoing collision, because tidal damages due to such a collision, particularly the removal of outer gas and stellar disk and the development of the characteristic tidal tails, should form only after the closest encounter. Physical separation by a direct collision, though naively attractive, is not a likely explanation because one cannot simply slice a gas cloud into two pieces. The face-on collision involving two disk galaxies in the ``Taffy'' pair resulted in only a slight re-arrangement and displacement of the HI in both disks \\citep{Condon93}, without producing any dramatic consequences of the type speculated here.  The general proximity and relative orientation of NGC~7320c with respect to the optical and HI tails have been invoked as evidence for its recent passage through SQ and possibly as an explanation for the northern optical tail that appears to originate from NGC~7319 \\citep[e.g.][]{Moles98,Sulentic01}. \\citet{Sulentic01} et al. even speculated that the two tidal tails might be the products of two consecutive passages by NGC~7320c. As clearly seen in various numerical studies of galaxy interactions \\citep[e.g.][]{BH96}, each passage by a companion generally erases all pre-existing features such as spiral arms and tidal tails, and such an explanation for the pair of tidal tails is probably not physical.  Another common tendency seen in numerical simulations is that tidal tails generally points back to the responsible culprit because the tidal debris and the passing companion have exchanged momentum, putting them in similar trajectories about the primary.  This means the northern HI tail, at least the bulk of the segment pointing due north, {\\it cannot be primarily driven by NGC~7320c}. The origin for the older (more diffuse and thicker) southern tail is also quite uncertain.  The number of faint neighbors with similar radial velocities as the galaxies in SQ is evidence that the quartet is not well isolated from the larger field of galaxies (see \\S~\\ref{sec:neighbors}).  Therefore interlopers may pass through and disturb the group with some frequency, including its HI companions such as Anon.~4.  Resulting dynamical heating may be sustaining the group from collapsing \\citep{Governato}, contrary to the prediction by \\citet{Barns96}.  One of the consequences of such a scenario is that interpreting the observed tidal features is no longer straightforward as in interacting galaxy pairs. The high galaxy density in core of the group also means that any galaxy passing through the group is subject to multiple scatterings rather than following a simple trajectory, and extrapolating the intuitions derived from the numerical studies of interacting pairs should be exercised with great caution.  An entirely different scenario one might consider, for at least parts of the observed HI features such as the large arc feature, is that these are remnants of the primordial gas clouds where the group has formed out of or tidal debris formed during the initial formation of the group, similar to the Leo ring found around the M96 group \\citep{Schneider85,Rood85,Schneider89}. The Leo Ring is a nearly complete ring of HI with a diameter of about 200 kpc, thought to be orbiting the galaxies M105 (E0) and NGC~3384 (SB0) with a period of $4\\times 10^9$ years. When examined at the same HI column density level of $9\\times 10^{19}$ cm$^{-2}$ shown in Figure~\\ref{fig:tot}, both the Leo Ring and the HI arc feature are partial rings of HI of $\\sim100$ kpc in size with linear velocity gradient perpendicular to their lengths. They both surround a group of galaxies interior to the ring, and their systemic velocities match the galaxy group velocity. One significant difference is that at least parts of the HI arc in SQ have stellar counterparts, which rejects the primordial origin.  The possibility that these rings of HI are remnants of group formation process, older than the group crossing time still seems plausible."
    },
    "0201/astro-ph0201345_arXiv.txt": {
        "abstract": "{ We show that the well-known discrepancy, known for about two decades, between the radial dependence of the Galactic cosmic ray nucleon distribution, as inferred most recently from EGRET observations of diffuse $\\gamma$-rays above 100 MeV, and of the most likely cosmic ray source distribution (supernova remnants, superbubbles, pulsars)  can be explained purely by {\\em propagation} effects. Contrary to previous claims, we demonstrate that this is possible, if the dynamical coupling between the escaping cosmic rays and the thermal plasma  is  taken into account, and thus a self-consistent calculation of a {\\em Galactic Wind} is carried out. Given a dependence of the cosmic ray source distribution on Galactocentric radius  $r$, our numerical wind solutions show that the cosmic ray outflow velocity, $V(r,z) = u_0 + V_{{\\rm A} 0}$, also depends both on $r$, as well as on  vertical distance $z$, with $u_0$ and $V_{{\\rm A} 0}$ denoting the thermal gas and the Alfv\\'en velocities, respectively, at a reference level $z_{\\rm C}$. The latter is by definition the transition boundary from diffusion to advection dominated cosmic ray transport and is therefore also a function of $r$. In fact,  the cosmic ray escape time averaged over particle energies decreases with increasing cosmic ray source strength. Thus an increase in cosmic ray source strength is counteracted by a reduced average cosmic ray residence time in the gas disk. This means that pronounced peaks in the radial distribution of the source strength result  in  mild radial $\\gamma$-ray gradients at GeV energies, as it has been observed. The effect might be enhanced by anisotropic diffusion, assuming different radial and vertical diffusion coefficients. In order to better understand the mechanism described, we have calculated analytic solutions of the stationary diffusion-advection equation, including anisotropic diffusion in an axisymmetric geometry, for a given cosmic ray source distribution and a realistic outflow velocity field $V(r,z)$, as inferred from the self-consistent numerical Galactic Wind simulations performed simultaneously. At TeV energies the $\\gamma$-rays from the sources themselves are expected to dominate the observed ``diffuse'' flux from the disk. Its observation should therefore allow an empirical test of the theory presented. ",
        "introduction": "Our information on the spatial distribution of cosmic rays (CRs) in the Galaxy stems largely from measurements of nonthermal emission, generated by the energetic charged particles interacting with matter and electromagnetic fields. For $\\gamma$-ray energies above 100 MeV, the main production process is probably via $\\pi^0$-decay, resulting from nuclear collisions between high energy particles and interstellar matter.  Past and recent observations in the GeV range have shown a roughly uniform distribution of diffuse $\\gamma$-ray emissivity in the Galactic plane, exhibiting only a shallow radial gradient (in a cylindrical Coordinate system). Hence, if $\\gamma$-rays were to map the spatial CR distribution, we would expect it to be uniform as well. However, associating CR production regions with star formation regions, all possible Galactic CR {\\em source} distributions are strongly peaked towards a Galactocentric distance at which a ring of molecular gas resides. It is commonly believed that the bulk of the CR nucleons below about $10^{15} \\, {\\rm eV}$ is produced in supernova remnants (SNRs) - the majority being core collapse SNRs - most likely by diffusive shock acceleration (Krymsky 1977; Axford et al. 1977; Bell 1978a,b;  Blandford  \\& Ostriker 1978). It has been  argued that up to some tenths of the hydrodynamic explosion energy might be converted  into CR energy  (e.g., Berezhko et al. 1994).  Since high-mass star formation mostly occurs in a spatially nonuniform manner, i.e.\\ in OB associations predominantly located in the spiral arms of late type galaxies, we are confronted with the problem of reproducing a mild radial gradient in the diffuse Galactic $\\gamma$-ray emission as it has been observed for the first time by the COS-B satellite (Strong et al. 1988) and, more recently, with higher angular and energy resolution, higher sensitivity and lower background by the EGRET instrument of the CGRO satellite (Strong \\& Mattox 1996, Digel et al. 1996). If the SNRs are the sources of the CR nucleon component and if the source distribution is inhomogeneous, this discrepancy must arise during the {\\em propagation} of CRs from their sources through the interstellar medium. Unlike e.g.\\ the interpretation of radio synchrotron emission, generated by relativistic electrons, $\\gamma$-ray data open the possibility of studying the nucleonic component of the CRs, in which almost all of the energy is stored (see e.g., Dogiel \\& Sch\\\"onfelder 1997). The distribution of the $\\gamma$-ray emissivity in the Galactic disk therefore bears important information on the  origin of CRs and on the conditions of CR propagation in the Galaxy. ",
        "conclusions": "\\label{discon} In the detailed analysis of the previous Sections it was shown that a physically reasonable CR transport model should include the nonlinear dynamical effects of the Galactic distribution of the CR sources themselves. Taking into account a simple radial dependence of the source power as suggested by radio observations of supernova remnants and pulsars, and thus averaging over all azimuthal angles, leads already to a radically different spatial behaviour of the CR distribution function. In all cases we have studied so far, there is a clear tendency to smooth out peaks due to enhanced particle injection by the sources in the ensuing radial particle distribution. This can be attributed to {\\em CR advection in an outflow} aided or even caused by the CRs themselves. In other words, the importance of CR advection is proportional to the CR power of the underlying sources. As a consequence, e.g.\\ the CR escape time, which is usually inferred from secondary radioactive isotopes, need no longer be a globally averaged quantity, but could vary strongly from place to place in the Galaxy. Therefore, the usual assumption that parameters of the CR flux measured near Earth are representative for the global processes of CR propagation in the Galaxy may be a misleading oversimplification.  It is common practice to determine from the intensities of different CR components near Earth the average diffusion coefficient in the Galaxy, the velocity of advection, the height of the CR halo in the direction perpendicular to the Galactic plane, and the CR injection spectrum, just to name the most important ones. Based on this hypothesis the nearly uniform radial CR distribution, derived from the measurement of diffuse Galactic $\\gamma$-rays, can be reproduced only if there exists thorough spatial mixing of CRs in the framework of an extended halo (if CR diffusion is isotropic). Hence in such a case {\\em local} and {\\em global} properties of CRs do not differ from each other. However, the inferred halo height from chemical composition ($z_{\\rm C} = 2 - 4$ kpc; see e.g., Bloemen et al.  1991, 1993; Webber et al. 1993; Lukasiak et al. 1994) is clearly inconsistent with the value derived from the interpretation of the $\\gamma$-ray data ($z_{\\rm C} = 15 - 20$ kpc; cf.\\ Dogiel \\& Uryson 1988, Bloemen et al.  1993) within the framework of an isotropic diffusion model (see Appendix~\\ref{appa}). We therefore conclude that the halo size derived from CR nuclear data reflects only a local value near Earth, and the huge halo extension derived previously from $\\gamma$-ray data may be an artifact, since it relies on the validity of global values for locally obtained CR data. This conclusion is supported by our numerical galactic wind simulations, which show that the vertical distance of the diffusion-advection transition boundary from the Galactic plane, is inversely proportional to the CR source power and not spatially constant as been previously assumed.  Radio observations of external galaxies indicate a large-scale magnetic field geometry, which is mainly parallel to the major axis in the disk, and if a halo field exists, it is parallel to the minor axis. Therefore we expect that CR diffusion is in general {\\em anisotropic}, with a radial diffusion coefficient $\\kappa_{\\rm r}$ in the disk, which is much larger than diffusion in the perpendicular direction, $D_{\\rm z}$, and vice versa in the halo. In this case the initially inhomogeneous CR distribution, due to a radially varying source distribution in the disk, is smeared out, whereas in the halo the dominant diffusion component $D_{\\rm z}$ can be superposed by a strong advection velocity, which may determine the spatial particle distribution.  It would be desirable to have a high enough spatial resolution and photon statistics in the future to observe the radial distribution of diffuse $\\gamma$-rays above 100 MeV in nearby edge-on galaxies, such as NGC253. However, it seems unlikely that both space-borne and ground-based $\\gamma$-ray observatories will satisfy this requirement in the near future. Thus the only direct observation of the CR source distribution in the Galaxy will be possible with next generation TeV instruments like H.E.S.S."
    },
    "0201/astro-ph0201035_arXiv.txt": {
        "abstract": "The structure of obscuring matter in the environment of active galactic nuclei with associated nuclear starbursts is investigated using 3-D hydrodynamical simulations. Simple analytical estimates suggest that the obscuring matter with energy feedback from supernovae has a torus-like structure with a radius of several tens of parsecs and a scale height of $\\sim 10$ pc. These estimates are confirmed by the fully non-linear numerical simulations, in which the multi-phase inhomogeneous interstellar matter and its interaction with the supernovae are consistently followed.  The globally stable, torus-like structure is highly inhomogeneous and turbulent. To achieve the high column densities ($ \\gtrsim 10^{24}$ cm$^{-2}$) as suggested by observations of some Seyfert 2 galaxies with nuclear starbursts, the viewing angle should be larger than about 70$^\\circ$ from the pole-on for a $10^8 M_\\odot$ massive black hole. Due to the inhomogeneous internal structure of the torus, the observed column density is sensitive to the line-of-sight, and it fluctuates by a factor of order $\\sim 100$. The covering fraction for $N > 10^{23}$ cm$^{-2}$ is about 0.4.  The average accretion rate toward $R < 1$ pc is $\\sim 0.4 M_\\odot$ yr$^{-1}$, which is boosted to twice that in the model without the energy feedback. ",
        "introduction": "Optically thick obscuring molecular tori have been postulated to explain various properties of active galactic nuclei (AGN), especially the two major categories of the AGN, namely type 1 and type 2.  However, the true structure and the formation and maintenance mechanism of the torus have not yet been understood either theoretically or observationally.  Recently a number of observations have suggested a new aspect of Seyfert 2 galaxies (Sy2).  It is pointed that there are possibly two types of Sy2s: one is a classic Sy2, an obscured Seyfert 1 nucleus. The other type II is a Sy2 with a nuclear starburst \\citep{cid95,gra97,gon01,lev01}.  In the latter case it is believed that a nuclear starburst is associated with the obscuring material whose scale is $R < 100 $ pc.  These observations imply that the classic picture of the unified model for AGNs, in which the diversity of the Seyfert galaxies is explained only by the geometrical effect of a geometrically and optically thick torus, may not be applicable to some fraction of the Sy2s (see also \\citet{maio95}).  More generally we propose here that there is an additional parameter in the unified models, namely, the strength of the nuclear starburst. It is this nuclear starburst and the mass of the black hole (BH) that determines the geometry of the obscuring torus.  While discussing the X-ray background, Fabian et al. (1998) proposed that low-luminosity AGN could be obscured by nuclear starbursts within the inner $\\sim 100$ pc.  They suggested that supernovae (SNe) from a nuclear starburst can provide the energy to boost the scale height of circumnuclear clouds and so obscure the nucleus. For simplicity, they assumed that the material sits in an isothermal sphere, and the gravitational effects of the central BH are ignored. This idea seems to be consistent with the recent observations of Sy2s with nuclear starbursts, Here we investigate the model in more detail.  The problem is actually very complicated. We need to solve the 3-D inhomogeneous structure and dynamics of the ISM in the combined gravitational potential of the central massive BH, the central stellar system and the central massive gas distribution.  We must handle the radiative cooling not only for the hot gas, but also for the cold gas, and include the feedback interaction between SNe and the inhomogeneous ISM.  Numerical simulations are powerful tools for this kind of complicated problem.  Recently we presented high-resolution, 3-D hydrodynamical modeling of the ISM in the central hundred pc region in galaxies, taking into account self-gravity of the gas, radiative cooling, and energy feedback from SNe \\citep{wad01b,wada01} (hereafter WN01 and W01). In this {\\it Letter}, we apply this numerical method to the gas dynamics around the central massive BH, and study the starburst-AGN connection in Seyfert galaxies. ",
        "conclusions": "The estimate in \\S 2 suggested that the opening angle positively correlates with $M_{\\rm BH}$. As $M_{\\rm BH}$ increases, the critical radius for the gravitational instability is also larger. This implies that low luminosity (i.e. small $M_{\\rm BH}$) AGNs are obscured for larger fraction of the viewing angle than the high luminosity AGN, if, quite reasonably, the mass accretion rate has a positive dependence on the mass of the central engine.  In other words, luminous type 2 objects, such as type 2 quasars are less frequently observed in the optical than type 1 quasars. They could be observed mainly X-rays, and this is consistent with observations \\citep{norman01,leh01,akiyama00}.  \\citet{kro88} (KB) claimed that stirring by stellar processes is never strong enough to compete with energy dissipation in the clumpy torus. One should note, however, that in their order of magnitude estimate, a much smaller, thicker torus ($\\sim 1$ pc) is assumed than in our extended disk model. Therefore KB need large velocity dispersions ($> 200$ km s$^{-1}$) to keep the disk thick, because thick disks with $h/r \\sim 1$ require $v_t/v_c \\sim 1$. The velocity dispersion for $h/r \\sim 1$ can be less than 50 km s$^{-1}$ for $r \\gtrsim 30$ pc. The smaller velocity dispersion is favorable, because the energy dissipation rate is smaller (see eq.(2)). KB also pointed out that a luminosity-to-gas-mass ratio, $L/M_g$, for the stellar stirring model would be a factor 100 larger than the observed values.  This is also not the case for a 100 pc scale obscuring material. For the analytical model presented here, $L/M_g \\sim r^{-3/2}$. In fact, a much larger molecular gas mass is expected on such a scale (e.g. \\citet{kohno99}). Recent IR and X-ray observations of Seyfert nuclei (e.g.  \\citet{gra97}), on the other hand, suggest a 100 pc scale extended obscuring material. This picture is consistent with our model of the extended starburst supported obscuring region in which case we still obtain a geometrically thick disk with velocity dispersions $\\lesssim 50$ km s$^{-1}$. Note also Pier \\& Krolik (1992) and Ohsuga \\& Umemura (2001) for the effects of radiation pressure on supporting the obscuring torus.  The present model suggests that the ratio of Sy2s to Sy1s cannot be explained only by the orientation of the viewer relative to the torus. In Sy1s, the star formation rate in the nuclear region would be smaller than in Sy2s.  The strict unified model, in this sense, should be modified by including the strength of the nuclear starburst."
    },
    "0201/astro-ph0201203_arXiv.txt": {
        "abstract": "The dynamics of accretion onto a schwarzschild black hole is studied using Paczynski-Wiita pseudo newtonian potential. Steady state solution of the flow equations is obtained using thin disc approximation, including the effect of bremsstrahlung cooling in the energy equation. The topology of transonic solutions are got and the conditions for shock formation (Rankine-Hugoniot conditions) are checked. Shock and no-shock regions in the parameter space of accretion rate and specific angular momentum are identified. We motivate the time-dependent numerical simulation of the flow, as a candidate for explaining the quasi periodic oscillations observed in black hole candidates. ",
        "introduction": "The parameter space of shock formation in adiabatic transonic accretion discs was first discussed in Chakrabarti(1989). Though time-dependent numerical simulation of transonic accretion disc in the presence of cooling was dealt in Molteni, Sponholz \\& Chakrabarti (1996, hereafter MSC96), the detailed study of parameter space was reserved for the future. In this paper we describe the procedure for obtaining the flow topology in the presence of bremsstrahlung cooling and obtain the dependence of topology and shock formation on the parameters, accretion rate and specific angular momentum. ",
        "conclusions": "Of all the different possible branches of accretion for chosen parameter value, the real flow is likely to choose the branch which is most stable, as perturbations are always present in a real situation. If the assumption of steady flow is relaxed the flow might still choose a steady solution branch if the time scale of change is large. Time-dependent numerical simulation of the flow (MSC96) shows that flow solution oscillates about the steady state solution when certain resonance condition is met. The 'perturbations' which occur in a numerical code and physical perturbations should be related, to increase the faith in numerical results. Such numerical studies will be pursued in the future. The oscillation of shock location (MSC96) would mean the size of post-shock region, which is the source for hard photons (Chakrabarti \\& Titarchuk 1995), is also oscillating. These would result in quasi periodic oscillations as reported in Chakrabarti \\& Manickam (2000)."
    },
    "0201/astro-ph0201291_arXiv.txt": {
        "abstract": "The measurements of cosmic interplanetary dust by the instruments on board the Pioneer 10 and 11 spacecraft contain the dynamical signature of dust generated by Edgeworth-Kuiper Belt objects, as well as short period Oort Cloud comets and short period Jupiter family comets. While the dust concentration detected between Jupiter and Saturn is mainly due to the cometary components, the dust outside Saturn's orbit is dominated by grains originating from the Edgeworth-Kuiper Belt. In order to sustain a dust concentration that accounts for the Pioneer measurements, short period external Jupiter family comets, on orbits similar to comet 29P/Schwassmann-Wachmann-1, have to produce $8\\times 10^4\\:{\\rm g}\\:{\\rm s}^{-1}$ of dust grains with sizes between $0.01$ and $6\\:{\\rm mm}$. A sustained production rate of $3\\times 10^5\\:{\\rm g}\\:{\\rm s}^{-1}$ has to be provided by short period Oort cloud comets on 1P/Halley-like orbits. The comets can not, however, account for the dust flux measured outside Saturn's orbit. The measurements there can only be explained by a generation of dust grains in the Edgeworth-Kuiper belt by mutual collisions of the source objects and by impacts of interstellar dust grains onto the objects' surfaces. These processes have to release in total $5\\times 10^7\\:{\\rm g}\\:{\\rm s}^{-1}$ of dust from the Edgeworth Kuiper belt objects in order to account for the amount of dust found by Pioneer beyond Saturn, making the Edgeworth-Kuiper disk the brightest extended feature of the Solar System when observed from afar. ",
        "introduction": "Our Solar System as well as other planetary systems is filled with small solid particles, either interstellar survivors of the formation process, or fragments of larger bodies like asteroids, comets, moons, or planets. Commonly referred to as interplanetary dust, these particles carry information about their sources, not only by their chemical signature \\citep{brownlee85,kissel86}, but also by the size and shape of their orbits around the Sun. The particles' chemistry as well as their orbit can best be gauged in situ, that is by dust detectors on board interplanetary spacecraft. While the accretion of interplanetary dust particles by the Earth's atmosphere allows their mineralogical, chemical, and isotopic analysis in ground-based laboratories after their collection by high-flying aircraft, information on their orbit around the Sun is lost after the atmospheric entry. The orbital properties of Solar System dust inside Jupiter's orbit has been extensively studied by in situ measurements \\citep{mcdonnell75,gruen77,gruen95b,gruen95c,brownlee97}. From these measurements Jupiter family short period comets and asteroids have been identified as the dominant dust sources \\citep{liou95,dermott92}. In the grain size regime below $1\\:{\\rm \\mu m}$ a high abundance of interstellar grains was found \\citep{gruen93}. While interstellar impactors can easily be distinguished from detections caused by solar system dust, it is still unclear what the relative contribution of the various interplanetary sources is. Besides this uncertainty, the large number of in situ measurements taken inside Jupiter's orbit led to a consistent picture of the extend and distribution solar system dust cloud there. The situation beyond Jupiter's orbit is however vastly different. So far the only in situ dust detectors ever to fly beyond Jupiter are the dust experiments on board the Pioneer 10 and 11 spacecraft \\citep{humes80}\\footnote{Since 31 December 2000, the Cassini spacecraft is outside Jupiter's orbit on its way to its final destination Saturn.}. Measurements of the plasma instruments on board Voyager 1 and 2 seem to indicate a high concentration of micron-sized particles out to $50\\:{\\rm AU}$ \\citep{gurnett97}. The Voyager results are however not conclusive because the plasma instruments have never been calibrated to measure dust impacts. From the Pioneer 10 and 11 measurements \\cite{humes80} found that, taken as an ensemble, the particles have to have a constant spatial concentration as function of the distance from the Sun and move on highly eccentric, randomly oriented orbits. In this report we use the Pioneer 10 and 11 data to identify the source objects of the particles by modelling the sources' signature in the Pioneer data, and comparing the measurements with the result of the modelling. ",
        "conclusions": "The discussion above shows that we have been able to identify a set of observable dust sources for the Pioneer dust measurements. Unlike the interpretation by \\cite{humes80}, we have used a set of 3 dynamical families of source objects. The sum of these sources provides the right spatial and local velocity distribution that explains the penetration fluxes measured by Pioneer. We found the calculated signature of the source families in the data to be independent, that is dominant for different heliocentric distances, so that dust production rates for the individual sources could be derived separately from the data. Especially the data collected by the spacecraft outside Saturn's orbit is very valuable, because with increasing heliocentric distance the number of possible contributors to the interplanetary dust cloud decreases. The only known source of interplanetary dust outside Saturn is the Edgeworth-Kuiper belt. This gives us the opportunity to unambiguously determine the amount of dust released by the objects of the belt. According to the Pioneer 10 measurements, the density of interplanetary dust generated by the Edgeworth-Kuiper belt is high enough so that this dust cloud is the second brightest feature of the solar system when observed from afar \\citep{liou99b}. Thus the Edgeworth-Kuiper belt and the distribution of dust particles it produces can act as a model for detecting other planetary systems around mid-age main sequence stars. Interplanetary dust in the region between Jupiter and Saturn gives us information about the dynamical properties of this interesting region. Since a fly-by of Jupiter on 31 December 2000 the Cassini spacecraft is on-route to Saturn, carrying a highly sensitive dust instrument. It will provide data on the mass, velocity, and chemical composition of the smaller sized dust particles."
    },
    "0201/astro-ph0201258_arXiv.txt": {
        "abstract": "{We compare the results of Balmer-line calculations using recent theory and improved computational algorithms with those from the widely-used SYNTHE and BALMER9 routines.  The resulting profiles are mostly indistinguishable. Good fits to the normalized solar Balmer lines H$\\alpha$ through H$\\delta$ are obtained (apart from the cores) using the recent unified-broadening calculations by Barklem and his coworkers provided that some adjustment for the continuum is performed. We discuss a surprising linearity with temperature of the Balmer line profiles in dwarfs. ",
        "introduction": "Balmer line strengths are highly sensitive to the temperature in cool stars because of the 10.2 eV excitation of the $n=2$ level from which they arise.  Fig. 151 from Uns\\\"{o}ld's (1955) classic text illustrates this for H$\\gamma$ equivalent widths. We show the effect in a different way in Fig.~\\ref{wing4.ps}, based on more recent line-broadening theory. The figure is for points on the H$\\alpha$ profile 4~\\AA\\ from the line center, but is characteristic of much of the line profile.   \\begin{figure} \\resizebox{9cm }{!} {\\includegraphics[0., 0.][450., 450.]{fig1.ps}} \\caption{H$\\alpha$ wing strength vs. $T_{\\rm eff}$ for several values of \\logg. The profiles are taken from the BP00K2NOVER grid available  in http://kurucz.harvard.edu} \\label{wing4.ps} \\end{figure}  An extensive investigation of Balmer lines in cool dwarfs (Fuhrmann, Axer, \\& Gehren 1993; Fuhrmann, Axer, \\& Gehren 1994) concluded these lines provide a more consistent guide to effective temperatures than broad-band colors or $b - y$.    Nevertheless, Balmer line profiles are not regularly used to fix the effective temperature of cool stars.  The reasons for this are numerous, but have not been explicitly addressed.  Some insight may be gained from the papers by Van't Veer-Menneret \\& M\\'{e}gessier (1996) or Castelli, Gratton \\& Kurucz (1997, henceforth, CGK).  A recent paper which does discuss use of H$\\alpha$ in the determination of effective temperatures is by Peterson, Dorman \\& Rood (2001). In addition to the uncertainties in placing the continuum level, uncertainties, both in the theory of stellar atmospheres ($l/H$, convection) and line formation remain unresolved.  The absorption coefficient of neutral hydrogen  takes into account the effects due to the natural absorption (natural broadening), the velocity of the absorbing hydrogen atoms (thermal Doppler and microturbulent broadening), the interactions with charged perturbers (linear Stark broadening),  with neutral perturbers different from hydrogen (van der Waals broadening), and with neutral hydrogen perturbers (resonance and van der Waals broadening). Each effect is represented by a profile and the total effect requires a convolution. Thermal Doppler and microturbulent broadenings are described by gaussian functions while natural, resonance, and van der Waals broadenings have  Lorentz profiles.   These two profiles are combined into a Voigt function.  The convolution of the Voigt profile with the Stark profile or Stark plus thermal Doppler effect then gives the  total absorption profile.  Most of the damping constants and Stark profiles are computed from complex theories based on several approximations, while the complete convolution of all the above profiles is a very time consuming algorithm.  In this paper we describe our attempts to evaluate several aspects of the calculations of Balmer line profiles.  \\section {Stark profiles}  Most work on stellar atmospheres makes use of codes provided by Kurucz (http://kurucz.harvard.edu). For computing hydrogen lines the codes are either BALMER9 (Kurucz, 1993a) which produces profiles for H$_{\\alpha}$, H$_{\\beta}$, H$_{\\gamma}$, and H$_{\\delta}$ or the SYNTHE code (Kurucz, 1993b) which produces profiles for any hydrogen line. In the first case  Stark profiles are interpolated in the Vidal, Cooper, \\& Smith (1973, henceforth VCS) tables, while in the second case the Stark profiles are based on the quasi-static Griem theory with parameters adjusted in such a way that profiles from Griem theory fit  the VCS profiles of the first members of the Lyman and Balmer series.  Only the most recent work on the Balmer lines (e.g. Barklem, Piskunov \\& O'Mara 2000, henceforth, BPO) has included the new Stark profiles of Chantal Stehl\\'{e} (henceforth CS) and her coworkers.  They are available from a link on her website: http://dasgal.obspm.fr/~stehle/. A recent reference is Stehl\\'{e} \\& Hutcheon (1999).  A problem arises when a given Stark profile is interpolated either in the VCS  or in the CS tables by using the interpolation method taken from the BALMER9 code. This is a bilinear interpolation in $\\log(T)$ and $\\log(N_e)$, followed by a linear interpolation in the parameter $\\Delta\\alpha = \\Delta\\lambda[{\\rm \\AA}]/F^0$.  Here, $F^0$ is the normal field strength in Gaussian cgs units, $F^0 = 1.25 N_e^{2/3}$, so the interpolation in $\\Delta\\alpha$ is not independent of the previous one which involves the electron density $N_e$. We find this introduces a small error that shows up as an oscillation in a plot of the Stark profile $S(\\Delta\\alpha)$ vs depth in the solar atmosphere for a small range of displacements from the line center as shown in Fig.~\\ref{waves2.ps}.  \\begin{figure} \\resizebox{8.0cm }{!} {\\includegraphics[angle=-90]{fig2.ps}} \\caption{Normalized Stark width at $\\Delta\\lambda = 0.5$\\AA \\ for H$\\alpha$ vs. 137 depths in an Holweger-M\\\"uller (1974) solar model. Each depth step is 0.05 in $\\log(\\tau_{\\lambda 5000})$. The vertical lines mark depths corresponding to boundaries of the tables giving $S(\\alpha)$ for a fixed value of the electron density. } \\label{waves2.ps} \\end{figure}  We were able to remove the oscillations by rewriting the CS tables with $\\Delta\\lambda$ as the third (independent) variable, and using essentially the same interpolation scheme as BALMER9.  Fortunately, it has resulted that the improved interpolation leads to no perceptible changes in the resulting line profiles. ",
        "conclusions": "We have explored recent techniques for computing Balmer line profiles in the sun, and H$\\alpha$ profiles in several models with temperatures ranging from 4500K to 12000K.  We find that new Stark profiles, rigorous convolution, and  improved  interpolation techniques make almost no difference in the resulting calculated profiles, compared with algorithms used in the Kurucz codes for several decades.  Good fits to normalized disk center solar profiles for the H$\\alpha$ through H$\\delta$ are obtained from the Holweger-M\\\"{u}ller (HM) model.  The H$\\alpha$ profile can also be reasonably fitted in absolute intensity, but the calculated continua for H$\\beta$ through H$\\delta$ are too high. This may reasonably be attributed to missing UV opacity, perhaps also to inadequacies of the HM model used here, as well as to uncertainties in the absolute solar calibration.  In spite of severe temperature inhomogeneities in the solar atmosphere, the plane-parallel model appears remarkably robust."
    },
    "0201/astro-ph0201544_arXiv.txt": {
        "abstract": "We present a comprehensive set of convergence tests which explore the role of various numerical parameters on the equilibrium structure of a simulated dark matter halo. We report results obtained with two independent, state-of-the-art, multi-stepping, parallel N--body codes: {\\tt PKDGRAV} and {\\tt GADGET}.  We find that convergent mass profiles can be obtained for suitable choices of the gravitational softening, timestep, force accuracy, initial redshift, and particle number. For softenings chosen so that particle discreteness effects are negligible, convergence in the circular velocity is obtained at radii where the following conditions are satisfied: (i) the timestep is much shorter than the local orbital timescale; (ii) accelerations do not exceed a characteristic acceleration imprinted by the gravitational softening; and (iii) enough particles are enclosed so that the collisional relaxation timescale is longer than the age of the universe.  Convergence also requires sufficiently high initial redshift and accurate force computations.  Poor spatial, time, or force resolution leads generally to systems with artificially low central density, but may also result in the formation of artificially dense central cusps.  We have explored several adaptive time-stepping choices and obtained best results when individual timesteps are chosen according to the local acceleration and the gravitational softening ($\\Delta t_i \\propto (\\epsilon/a_i)^{1/2}$), although further experimentation may yield better and more efficient criteria. The most stringent requirement for convergence is typically that imposed on the particle number by the collisional relaxation criterion, which implies that in order to estimate accurate circular velocities at radii where the density contrast may reach $\\sim 10^6$, the region must enclose of order $3000$ particles (or more than a few times $10^6$ within the virial radius).  Applying these criteria to a galaxy-sized $\\Lambda$CDM halo, we find that the spherically-averaged density profile becomes progressively shallower from the virial radius inwards, reaching a logarithmic slope shallower than $-1.2$ at the innermost resolved point, $r \\sim 0.005 \\, r_{200}$, with little evidence for convergence to a power-law behaviour in the inner regions. ",
        "introduction": "\\label{sec:intro} Over the past few decades, cosmological N-body simulations have led to impressive strides in our understanding of structure formation in universes dominated by collisionless dark matter. Such simulations have provided an ideal test-bed for analytic theories of structure formation, and have been used to validate and motivate a variety of theoretical insights into the statistics of hierarchical clustering (e.g., Press \\& Schechter 1974, Bardeen et al. 1986, Bond et al. 1991, Lacey \\& Cole 1993, Mo \\& White 1996). In particular, N-body simulations have played a pivotal role in providing a clear framework within which the CDM cosmogony may be compared with observation, and in establishing Cold Dark Matter (CDM) as the leading theory of structure formation (Davis et al. 1985).  This work has led to the development of a robust theoretical framework which provides an accurate statistical description of structure growth through gravitational instability seeded by Gaussian primordial density fluctuations. It is now possible to predict with great accuracy, and based only on the initial power spectrum of the primordial fluctuations, a number of important statistics that characterize the large scale structure of the universe; e.g., the mass function and clustering of dark matter halos and their evolution with redshift (e.g., Jing 1998, Sheth \\& Tormen 1999, Jenkins et al. 2001) the non-linear evolution of the dark matter power spectrum and correlation functions (e.g., Hamilton et al. 1991, Peacock \\& Dodds 1996), as well as the topological properties of the large scale structure (e.g., Gott, Weinberg \\& Melott 1987).  The impact of such simulation work has been greatest in the non-linear regime, where analytic calculations offer little guidance.  Recently, and as a result of the development of efficient algorithms and of the advent of massively parallel computers, it has been possible to apply N-body studies to the investigation of structure on small, highly non-linear scales. These studies can now probe scales comparable to the luminous radii of individual galaxies, thus enabling direct comparison between theory and observation in regions where luminous dynamical tracers are abundant and easiest to observe.  Predicting the structure of dark matter halos on kpc and sub-kpc scales, where it can be compared directly with observations of galactic dynamics, is one of the premier goals of N-body experiments, and there has been steady progress in this area over the past few years.  Building upon the early work of Frenk et al. (1985, 1988), Quinn, Salmon \\& Zurek (1986), Dubinski \\& Carlberg (1991) and Crone, Evrard \\& Richstone (1993), Navarro, Frenk \\& White (1996, 1997, hereafter NFW) found that, independently of mass and of the value of the cosmological parameters, the density profiles of dark matter halos formed in various hierarchical clustering cosmogonies were strikingly similar. This `universal' structure can be characterized by a spherically-averaged density profile which differs substantially from the simple power law, $\\rho(r) \\propto r^{-\\beta}$, predicted by early theoretical studies (Gunn \\& Gott 1972, Fillmore \\& Goldreich 1984, Hoffmann \\& Shaham 1985, White \\& Zaritsky 1992). The profile steepens monotonically with radius, with logarithmic slopes shallower than isothermal (i.e. $\\beta < 2$) near the centre, but steeper than isothermal ($\\beta>2$) in the outer regions.  NFW proposed a simple formula, \\begin{equation} \\label{eq:nfw} {\\rho(r) \\over \\rho_{\\rm crit}} = {\\delta_c \\over (r/r_s)(1+r/r_s)^2}, \\end{equation} which describes the density profile of any halo with only two parameters, a characteristic density contrast{\\footnote{We use the term `density contrast' to denote densities expressed in units of the critical density for closure, $\\rho_{\\rm crit}=3H^2/8\\pi G$. We express the present value of Hubble's constant as $H(z=0)=H_0=100\\, h$ km s$^{-1}$ Mpc$^{-1}$}}, $\\delta_c$, and a scale radius, $r_s$. Defining the mass of a halo as that contained within $r_{200}$, the radius of a sphere of mean density contrast $200$, there is a single adjustable parameter that fully describes the mass profile of halos of given mass: the `concentration' ratio $c=r_{200}/r_s$.  For the sake of this discussion, the two main points to note from the work of NFW are the following: (i) the density profile in the inner regions of the halo is shallower, and in the outer regions steeper, than isothermal, and (ii) there is no well defined value for the {\\it central} density of the dark matter, which can in principle climb to arbitrarily large values near the centre.  Conclusion (i) is important, since it is a feature of dark halo models that is required by observations. For example, it implies that the characteristic speeds of dynamical tracers may be lower near the centre than in the main body of the system, as observed in disk galaxies, where the velocity dispersion of the bulge is lower than indicated by the maximum rotation speed of the surrounding disk, as well as in galaxy clusters, where the velocity dispersion of stars in the central cluster galaxy is lower than that of the cluster as a whole. Conclusion (ii) is also important, since there have been a number of reports in the literature arguing that the shape of the rotation curves of many disk galaxies rules out steeply divergent dark matter density profiles (Flores \\& Primack 1994, Moore 1994, de Blok et al. 2001, but see van den Bosch \\& Swaters 2001), a result that may signal a genuine crisis for the CDM paradigm on small scales (see, e.g., Sellwood \\& Kosowsky 2000, Moore 2001).  These general results of the work by NFW have been confirmed by a number of subsequent studies (Cole \\& Lacey 1996, Fukushige \\& Makino 1997, Huss, Jain \\& Steinmetz 1999, Moore et al. 1998, Jing \\& Suto 2000), although there is some disagreement about the innermost value of the logarithmic slope. Moore et al. (1998), Ghigna et al.(2000), and Fukushige \\& Makino (1997, 2001) have argued that density profiles diverge near the centre with logarithmic slopes considerably steeper than the asymptotic value of $\\beta=1$ in NFW's formula. Kravtsov et al. (1998), on the other hand, initially obtained much shallower inner slopes ($\\beta \\sim 0.7$) in their numerical simulations, but have now revised their conclusions; these authors now argue that CDM halos have steeply divergent density profiles but, depending on evolutionary details, the slope of a galaxy-sized halo at the innermost resolved radius may vary between $-1.0$ and $-1.5$ (Klypin et al. 2001).  Since steep inner slopes are apparently disfavoured by rotation curve data it is important to establish this result conclusively; if confirmed, it may offer a way to falsify the CDM paradigm on small scales. Unfortunately, observational constraints are strongest just where theoretical predictions are least trustworthy. For example, the alleged disagreement between observed rotation curves and cuspy dark halo models is most evident for sub-$L_{\\star}$ galaxies on scales of $\\sim 1\\, h^{-1}$ kpc or less. For typical circular speeds of $\\sim 100$ km s$^{-1}$, this corresponds to regions where the density contrast exceeds $\\sim 10^6$. Orbital times in these regions are of order $10^{-3}$ of the age of the universe, implying that N-body codes must be able to follow particles accurately for several thousand orbits. Few cosmological codes have been tested in a systematic way under such circumstances. Furthermore, the cold dark matter halos that host typical disk galaxies are thought to extend out to a few hundred kpc, implying that the $\\sim$kpc scale probed by observations involves a very small fraction of the mass and volume of the dark halo. As a consequence, these regions are vulnerable to numerical artifacts in N-body simulations stemming, for example, from the gravitational softening or the number of particles.  Extreme care is thus needed to separate numerical artifacts from the true predictions of the Cold Dark Matter model. In order to validate or `rule out' the CDM cosmogony one must be certain that model predictions on the relevant scales are accurate, robust, and free of systematic numerical uncertainties. Although there have been some recent attempts at unravelling the role of numerical parameters on the structure of simulated dark matter halos, notably in the work of Moore et al. (1998), Knebe et al. (2000), Klypin et al. (2001) and Ghigna et al. (2000), the conclusions from these works are still preliminary and, in some cases, even contradictory.  To cite an example, Moore et al. (1998) argue that the smallest resolved scales correspond to about half the mean inter-particle separation within the virial radius, and conclude that many thousands of particles are needed to resolve the inner density profile of dark matter halos. Klypin et al. (2001), on the other hand, conclude that mass profiles can always be trusted down to the scale of the innermost $\\sim 200$ particles, provided that other numerical parameters are chosen wisely. Ghigna et al. (2000) suggest an additional convergence criterion based on the gravitational softening length scale, and argue that convergence is only achieved on scales that contain many particles and that are larger than about $\\sim 3$ times the scale where pairwise forces become Newtonian. Understanding the origin of such disparate conclusions and the precise role of numerical parameters is clearly needed before a firm theoretical prediction for the structure of CDM halos on $\\sim$kpc scales may emerge.  Motivated by this, we have undertaken a large series of numerical simulations designed to clarify the role of numerical parameters on the structure of simulated cold dark matter halos. In particular, we would like to answer the following question: what regions of a simulated dark matter halo in virial equilibrium can be considered reliably resolved?  This question is particularly difficult because of the lack of a theory with which the true structure of dark halos may be predicted analytically, so the best we can do is to establish the conditions under which the structure of a simulated dark halo is independent of numerical parameters. This is the question which we endeavor to answer in this paper.  There is a long list of considerations and numerical parameters that may influence the structure of simulated dark halos: \\begin{itemize} \\item{the N-body code itself} \\item{the procedure for generating initial conditions} \\item{the accuracy of the force computation} \\item{the integration scheme} \\item{the initial redshift} \\item{the time-stepping choice} \\item{the gravitational softening} \\item{the particle number} \\end{itemize} Clearly the list could be substantially longer, but the items above are widely considered the most important concerning the structure of simulated dark halos. Before we proceed to analyze their role, we must decide which properties of a dark matter halo we will assess for numerical convergence. Because, as mentioned above, disk galaxy rotation curves seem to pose at present one of the most pressing challenges to the CDM paradigm on small scales, we have decided to concentrate on the spherically-averaged mass profile, as measured by the radial dependence of the circular velocity, $V_c(r)=\\sqrt{GM(r)/r}$, or, equivalently, by the inner mean density profile, ${\\bar \\rho}(r)=3M(r)/4\\pi r^3$.  We note that the convergence criteria derived here apply strictly only to these properties, and that others, such as the three-dimensional shape of halos, their detailed orbital structure, or the mass function of substructure halos, may require different convergence criteria.  The basic philosophy of our convergence testing procedure is to select a small sample of halos from a cosmological simulation of a large periodic box and to resimulate them varying systematically the parameters listed above, searching for regions in parameter space where the circular velocity curves are independent of the value of the numerical parameters, down to the smallest scales where Poisson uncertainties become important, i.e., roughly down to the radius that contains $\\sim 100$ particles.  Overall, this is a fairly technical paper of interest mostly to practitioners of cosmological N-body simulations. Readers less interested in numerical details may wish to skip to \\S~\\ref{sec:vprof}, where we discuss in detail the converged inner mass profile of the galaxy-sized $\\Lambda$CDM halo used in our convergence study. The more technical sections include: \\begin{itemize} \\item {a discussion of the N-body codes used in this work, initial conditions setup, and analysis procedure (\\S~\\ref{sec:method} and Appendix)} \\item {a general discussion of the consequences of discreteness effects on simulations of dark matter halos, including a derivation of ``optimal'' choices (for given particle number) of the timestep and the gravitational softening (\\S~\\ref{sec:nepsdt})} \\item {a comparison between single- and multi-timestepping techniques (\\S~\\ref{sec:multistep}) } \\item {a discussion of the role of the gravitational softening, the initial redshift, the force accuracy, and the particle number on the inner mass profile of simulated halos (\\S~\\ref{sec:numpar}) } \\end{itemize} Finally, a worked example of how to choose optimal parameters for a high-resolution simulation is presented in \\S~\\ref{ssec:workex}. We summarize our main conclusions in \\S~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions} We have performed a comprehensive series of convergence tests designed to study the effect of numerical parameters on the structure of simulated cold dark matter halos. Our tests explore the influence of the gravitational softening, the time-stepping algorithm, the starting redshift, the accuracy of force computations, and the number of particles on the spherically-averaged mass profile of a galaxy-sized halo in the $\\Lambda$CDM cosmogony. We derive, for each of these parameters, empirical rules that optimize their choice or, when those choices are dictated by computational limitations, we offer simple prescriptions to assess the effective convergence of the mass profile of a simulated halo. Our main results can be summarized as follows:  \\begin{enumerate}  \\item {\\it Timestep and Discreteness Effects.}  The number of timesteps required to achieve convergence depends primarily on the orbital timescale of the region to be resolved, but may also be sensitive to the number of particles and the gravitational softening, unless these parameters are chosen so that discreteness effects are unimportant. This requires the gravitational softening to be large enough so that the maximum acceleration during two-body encounters does not exceed the minimum mean field acceleration in the halo, $\\epsilon \\gsim \\epsilon_{\\rm acc}=r_{200}/\\sqrt{N_{200}}$. Empirically, we find that $\\epsilon\\approx \\epsilon_{\\rm opt}=4 \\, \\epsilon_{\\rm acc}$ gives good results. When this condition is satisfied, the minimum converged radius, $r_{\\rm conv}$, is given by the condition that the circular orbit timescale should be long compared to the timestep, $t_{\\rm circ}(r_{\\rm conv}) \\approx 15 \\left({\\Delta t/ t_0}\\right)^{5/6} t_{\\rm circ}(r_{200})$.  {\\it Substantially smaller} timesteps are needed if $\\epsilon<\\epsilon_{\\rm opt}$. Dark matter densities at $r< r_{\\rm conv}$ may be under- or over-estimated, depending on the integrator and timestepping schemes used. For example, constant-timestep {\\tt GADGET} runs develop artificially dense, `cuspy' cores in poorly resolved regions, indicating that the approach to convergence is not always monotonic. This emphasizes the importance of comprehensive convergence tests such as the ones presented here to validate the results of numerical studies of the inner structure of CDM halos.  \\item {\\it Fixed Timestep versus Adaptive Multi-Stepping.}  Of the several adaptive, multiple time-stepping criteria that we considered, we have found best results when timesteps are chosen to depend explicitly on the gravitational softening and on the acceleration, $\\Delta t_i=\\eta_{a\\hskip-0.5pt\\epsilon}\\sqrt{\\epsilon_i/a_i}$, with $\\eta_{a\\hskip-0.5pt\\epsilon} \\sim 0.2$. Experiments with time-stepping choices that do not include explicitly the gravitational softening require the value of the corresponding $\\eta$ to be reduced as $\\epsilon$ is reduced below the optimal value in order to obtain convergence. In terms of computational cost, we find that multi-time-stepping criteria significantly outperform the use of a single timestep for all particles only for softenings well below the optimal value.  \\item {\\it Gravitational Softening.} The choice of gravitational softening is found to impose a maximum acceleration scale above which simulation results cannot be trusted. This acceleration scale appears to depend mainly on the circular velocity of the halo and on the gravitational softening scale, and is given by $a_{\\epsilon}=\\chi_{\\epsilon}\\, V_{200}^2/\\epsilon$, with $\\chi_{\\epsilon} \\sim 0.5$. For {\\it given particle number}, convergence to better than $10\\%$ in the mass profile is obtained at radii greater than $\\epsilon$ that also contain more than $100$ particles and where the acceleration criterion is satisfied: $a(r)=V_c(r)^2/r \\lsim a_{\\epsilon}$.  \\item {\\it Starting Redshift.} The mass profiles of simulated dark halos converge provided that the initial redshift is chosen so that the theoretical (linear) rms fluctuations on the smallest resolved mass scale, $m_p$ (the mass of one high-resolution particle) is $\\sigma(m_p,z_i) \\lsim 0.3$. Since $\\sigma(m_p)$ is a weak function of mass on subgalactic mass scales for CDM-like power spectra, this criterion indicates that a modest starting redshift, such as $1+z_i\\approx 50$ is appropriate for particle masses as low as $m_p \\sim 10^5 \\, h^{-1} M_{\\odot}$ in the $\\Lambda$CDM cosmogony.  \\item {\\it Force Accuracy.} The mass profiles of simulated CDM halos are quite sensitive to the accuracy of the force calculations, and convergence requires care in the choice of node opening criteria in the treecodes used in our study. Poor force accuracy leads to the development of artificially low density cores.  In the case of {\\tt GADGET}, for example, we find that even occasional large errors in the forces may lead to noticeable deviations from converged profiles. To avoid this, it is necessary to choose tree-walking parameters that curtail drastically the tail of the most deviant force calculations, however rare. In {\\tt GADGET} this can be achieved by activating the compiler option {\\tt -DBMAX}. Using up to hexadecapole terms in the node potential expansion and setting a redshift dependent tree-node opening criterion, as in {\\tt PKDGRAV}, where $\\theta=0.55$ is chosen for $z>2$ and $\\theta=0.7$ for $z<2$, seems also to work well.  \\item {\\it Particle Number.}  In order to achieve convergence in the mass profile, enough particles must be enclosed so that the average two-body relaxation timescale within the region is comparable or longer than the age of the universe. We find empirically that the condition, $t_{\\rm relax}(r)\\gsim 0.6 \\, t_0$, describes converged regions well. Since $t_{\\rm relax}$ is roughly proportional to the enclosed number of particles times the local dynamical timescale, resolving regions near the centre, where density contrasts are high and dynamical timescales are short, requires substantially more particles than resolving regions more distant from the centre. Of order $3000$ enclosed particles are needed to resolve regions where the density contrast reaches $10^6$. On the other hand, density contrasts of order $10^{4.5}$ require only $100$ enclosed particles for numerical convergence. Resolving radii of order $0.5\\%$ of the virial radius in the first case requires of order $3 \\times 10^6$ particles within the virial radius.  \\end{enumerate}  For most simulations, the most stringent convergence criterion is the relaxation timescale condition on the number of particles. This implies that there is little choice but to strive for the largest possible $N$ when studying the inner regions of dark matter halos. This limit is dictated by the available computer resources. Choosing the optimal softening for the adopted number of particles then minimizes the number of timesteps needed to achieve convergence down to the radius where $t_{\\rm relax}(r)\\gsim 0.6 \\, t_0$. The precise number of timesteps cannot be determined ahead of time, since $t_{\\rm relax}(r)$ depends on the detailed structure of the halo, which is what we are trying to measure. This implies that a series of simulations where the number of particles is increased gradually is advisable in order to ensure that optimal parameters are chosen for the highest-resolution run intended.  We have applied our convergence criteria to a $\\sim 205$ km s$^{-1}$ $\\Lambda$CDM halo in order to investigate the behaviour of the inner slope of the density profile. We find that the slope of the spherically-averaged density profile, $\\beta=-{\\rm d}\\log({\\rho})/{\\rm d}\\log(r)$, becomes increasingly shallow inwards, with little sign of approach to an asymptotic value.  At the smallest radius that we consider resolved in our highest-resolution ($256^3$) simulation ($r_{\\rm min} \\sim 1 \\, h^{-1}$ kpc $\\approx 0.005\\, r_{200}$), the local and cumulative density contrasts are robustly determined, $\\rho(r_{\\rm min})/\\rho_{\\rm crit}=9.4\\times 10^5$, and ${\\bar \\rho}(r_{\\rm min})/\\rho_{\\rm crit}\\approx 1.6\\times 10^6$. These values can be combined with the requirement of mass conservation to place an upper limit to the inner asymptotic slope of the density profile, $\\beta < 3 \\, (1-\\rho(r_{\\rm min})/{\\bar \\rho(r_{\\rm min})})=1.2$, although it is possible that the slope may actually become even shallower near the centre, as suggested recently by Taylor \\& Navarro (2001).  Our results thus argue against the very steep values for the asymptotic central slope ($\\beta\\approx 1.5$) claimed recently by Moore et al. (1998, 1999), Ghigna et al. (1998, 2000), and Fukushige and Makino (1997, 2001). The reasons for this disagreement are unclear at this point, since there are substantial differences in the halo mass, numerical techniques, and cosmological model adopted, which hinder a direct comparison between our results and theirs. For example, the work of Moore et al. (1998) and Ghigna et al. (2000) differs from ours in mass scale (they simulated a galaxy cluster while we target a galaxy-sized halo) and in cosmology (they adopted an Einstein-de Sitter CDM cosmogony, whereas we adopt the $\\Lambda$CDM model).  Finally, the difference between the conclusions from various authors may just reflect the fact that each group applies different criteria to the identification of the regions deemed trustworthy. We note that models with the very steep ($\\beta\\sim 1.5$) inner slopes proposed by the Moore et al group and with the shallower slopes that we find here are almost indistinguishable if we restrict our analysis to radii $\\gsim 2\\%$ of the virial radius. Probing radii within the inner $1\\%$ of the virial radius seems required to shed light on this controversy. Further simulation work with resolution adequate to address this issue in detail is currently underway."
    },
    "0201/astro-ph0201402_arXiv.txt": {
        "abstract": "The existence of the Fundamental Plane imposes strong constraints on the structure and dynamics of elliptical galaxies, and thus contains important information on the processes of their formation and evolution.  Here we focus on the relations between the Fundamental Plane thinness and tilt and the amount of radial orbital anisotropy: in fact, the problem of the compatibility between the observed thinness of the Fundamental Plane and the wide spread of orbital anisotropy admitted by galaxy models has been often risen. By using N-body simulations of galaxy models characterized by observationally motivated density profiles, and also allowing for the presence of live, massive dark matter halos, we explore the impact of radial orbital anisotropy and instability on the Fundamental Plane properties.  The numerical results confirm a previous semi--analytical finding (based on a different class of one--component galaxy models): the requirement of stability matches almost exactly the thinness of the Fundamental Plane.  In other words, galaxy models that are radially anisotropic enough to be found outside the observed Fundamental Plane (with their isotropic parent models lying on the Fundamental Plane) are unstable, and their end--products fall back on the Fundamental Plane itself.  We also find that a systematic increase of radial orbit anisotropy with galaxy luminosity cannot explain by itself the whole tilt of the Fundamental Plane, becoming the galaxy models unstable at moderately high luminosities: at variance with the previous case their end--products are found well outside the Fundamental Plane itself. Some physical implications of these findings are discussed in detail. ",
        "introduction": "The Fundamental Plane (FP) of elliptical galaxies (Djorgovsky \\& Davies 1987; Dressler et al. 1987) is a scaling relation between three of their basic {\\it observational} properties, namely the circularized effective radius $\\cRe\\equiv\\sqrt{\\sae\\sbe}$ (where $\\sae$ and $\\sbe$ are the major and minor semi-axis of the effective isophotal ellipse), the central velocity dispersion $\\sg0$, and the mean effective surface brightness $\\Ie{\\equiv}\\Lb/2\\pi{\\cRe}^2$ (where $\\Lb$ is the luminosity of the galaxy, for example in the Johnson B-band).  An interesting parameterization of the FP, that we adopt in this paper, has been introduced by Bender, Burstein \\& Faber (1992, hereafter BBF): \\begin{equation} \\ku\\equiv\\frac{\\log\\sg0^2 + \\log\\cRe}{\\sqrt 2}, \\end{equation} \\begin{equation} \\kd\\equiv\\frac{\\log\\sg0^2 + 2\\log\\Ie - \\log\\cRe}{\\sqrt 6}, \\end{equation} \\begin{equation} \\kt\\equiv\\frac{\\log\\sg0^2 - \\log\\Ie -\\log\\cRe}{\\sqrt 3}. \\end{equation} In particular, when projected on the $(\\ku ,\\kt)$ plane, the FP is seen almost edge--on and it is considerably thin, while the distribution of galaxies in the $(\\ku ,\\kd)$ plane is considerably broader.  For example, Virgo ellipticals studied by BBF are distributed on the $(\\ku ,\\kt)$ plane according to the best--fit relation \\begin{equation} \\kt=0.15\\ku+0.36 \\end{equation} (when adopting respectively, kpc, km s$^{-1}$ and $\\Lbsol$ pc$^{-2}$ as length, velocity and surface brightness units), with a very small dispersion of $\\sigma (\\kt)\\simeq 0.05$ over all the range spanned by the data, $2.6\\, \\lsim\\, \\ku\\, \\lsim\\, 4.6$ (and so $0.75\\, \\lsim \\, \\kt\\, \\lsim \\, 1.05$, see, e.g., Ciotti, Lanzoni \\& Renzini 1996, hereafter CLR96).  By combining equations (1) and (3) with equation (4) the FP equation of BBF is then obtained directly in terms of the observables: the exponents are in good agreement with those derived (in the Johnson B-band) from a much larger galaxy sample by J{\\o}rgensen, Franx \\& Kj{\\ae}rgaard (1996).  Note that equation (4) implies that for galaxies of given luminosity $\\Lb$ their effective radius and central velocity dispersion must be strongly coupled: in fact at any  fixed luminosity the coordinates $\\kt$ and $\\ku$ are related through definitions (1) and (3) by \\begin{equation} \\kt=\\sqrt{\\frac{2}{3}}\\ku+{\\sqrt{\\frac{1}{3}}}\\log {\\frac{2\\pi}{\\Lb}}, \\end{equation} and the slope of this relation ($\\simeq 0.82$) is different from that of the FP ($\\simeq 0.15$). As a consequence, in the $(\\ku,\\kt)$ plane all galaxies with the same luminosity are located on straight lines significantly inclined with respect to the FP: the presence of substantial scatter in galactic properties from galaxy to galaxy (of similar luminosity) would destroy the thinness of the FP by producing a large scatter in $\\ku$ and so in $\\kt$.  The relation between galaxy properties and the FP can be expressed in a quantitative way under the reasonable assumption that present day ellipticals are virialized systems. We write the virial theorem as \\begin{equation} \\frac{G\\Lb\\ml}{\\cRe}=\\Kv\\sg0^2 \\end{equation} where $\\ml=\\Mstar/\\Lb$ is the galaxy {\\it stellar} mass--to--light ratio (for example in Blue solar units), and $\\Kv$ is a dimensionless factor depending on the stellar density profile, internal dynamics, dark matter amount and distribution, and, for non--spherical galaxies, on their relative orientation with respect to the observer's line--of--sight (see, e.g., Ciotti 1997); in addition, $\\Kv$ depends also on the observing aperture adopted to derive $\\sg0$.  Equations (4) and (6) imply that, for galaxies belonging to the FP, the quantity $\\ml/\\Kv$ is a very well defined function of two\\footnote{In principle one could find virialized galaxies {\\it everywhere} in the three--dimensional observational ($\\Lb, \\sg0, \\cRe$) space.  From this point of view the existence of the FP is related to the virial theorem as the HR diagram is related to hydrostatic equilibrium: the useful information that one derives is {\\it not} about the equilibrium equations, but the physics of the objects involved.} of the three observables $\\Lb$, $\\sg0$, and $\\cRe$. In the particular case of adopting the numerical coefficients of equation (4), measuring $\\Lb$ in $10^{10}\\Lbsol$, and taking into account that $\\ku=\\log (G\\ml\\Lb/\\Kv)/\\sqrt{2}$ and $\\kt=\\log (2\\pi G\\ml/\\Kv)/\\sqrt{3}$, \\begin{equation} \\frac{\\ml}{\\Kv}\\simeq 1.12\\times{\\Lb}^{0.23} \\end{equation} where the quantity $\\ml/\\Kv$ is characterized by a scatter of $\\simeq 20\\%$ due to its relation with $\\kt$ (however, additional considerations reduce this figure to $\\simeq 12\\%$, see Renzini \\& Ciotti 1993).  As a consequence, any departure from the relation dictated by equation (7) will move a galaxy away from the FP.  We recall here that the dependence of $\\kt$ on $\\ku$, as given by equation (4) and responsible for the luminosity dependence of the ratio $\\ml/\\Kv$, is commonly known as the FP ``tilt''.  The simple analysis presented above shows that the two properties of thinness and tilt of the FP are deeply connected with the present--day structure and dynamics of ellipticals (hereafter, Es), and, as a consequence, with their formation and evolution history: the very existence of the FP [as well as of the other tight scaling relations revealing the remarkable homogeneity of Es, such as the $M_{\\rm BH}-\\sg0$ (Gebhardt et al.  2000; Ferrarese \\& Merritt 2000), the ${\\rm Mg}_2$-$\\sg0$ (Bender, Burstein \\& Faber, 1993 and references therein) and the color--magnitude (Bower, Lucey \\& Ellis, 1992) relations] imposes strong constraints on the different formation and evolutionary scenarios proposed for Es (i.e., dissipationless merging, monolithic collapse, or a combination of the two; see, e.g., Ciotti \\& van Albada 2001).  Among the various galaxy properties in principle able to destroy the FP thinness (as a consequence of a substantial variation at fixed galaxy luminosity), one of the most ``effective'' is certainly orbital anisotropy (de Zeeuw \\& Franx 1991). In fact, galaxy models are commonly believed to be able to sustain a large spread of orbital anisotropies and it is also well known that radial orbital anisotropy can produce very high {\\it central} velocity dispersion values, and correspondingly low values of $\\Kv$, thus substantially violating equation (7).  A natural question to be addressed is then what physical principle or evolutionary process limits the range of orbital anisotropies shown by real galaxies.  Ciotti \\& Lanzoni (1997, hereafter CL97), using one--component, radially anisotropic Sersic (1968) models, and a semi--analytical investigation based on the Fridman \\& Polyachenko (1984, hereafter FP84) stability indicator, suggested the possibility that radial orbit instability could be the limiting factor of the FP thickness.  In practice, CL97 found indications that galaxy models sufficiently anisotropic to be outside the FP observed thickness (when their parent isotropic model was assumed to lie on the FP) were unstable.  Clearly, this preliminary indication requires a confirmation with the aid of numerical simulations and more realistic galaxy models (for example, allowing for the presence of live dark matter halos). Also, a question naturally associated with that above is the determination of the position of the end--products of unstable initial conditions in the space of the observables.  One of the aims of this work is indeed to answer these two questions by numerical simulations of one and two--component radially anisotropic galaxy models.  Orbital anisotropy is not only related to the problem of the FP thinness but it is also one of the candidates that have been proposed to explain the origin of the FP tilt (CLR96; CL97). If this were the case, the amount of radial anisotropy in the velocity dispersion tensor should increase with galaxy luminosity, as can be seen from equation (7) under the assumption of a constant $\\ml$ and of structural homology\\footnote{With ``structural homology'' we mean that all the structural galaxy properties (e.g., the stellar and dark halo density profiles, the ratio of their scale--lengths and masses, and so on) {\\it do not depend} on $\\Lb$.} over the whole FP plane.  In other words, in this scenario the FP tilt would be produced by a {\\it dynamical} non--homology due to anisotropy (note that dynamical non--homology may well coexist with structural homology, but the converse is in general not true). Many interesting questions are raised by the scenario depicted above: for example, under the assumption that an isotropic galaxy of given luminosity lies on the FP, how far can the derived structurally homologous but radially anisotropic models climb over the FP before the onset of radial orbit instability?  In addition, what happens to the end--products of the unstable models?  Will they remain near the FP?  In this paper we try to address also these questions with the aid of N-body numerical simulations.  Strictly related to the clarification of the interplay between orbital anisotropy and the FP tilt, is the possibility to obtain some clues on the formation processes of Es.  In fact it is trivial to prove that, as consequence of the virial theorem and the conservation of the total energy, in the merging of two galaxies with masses $M_1$ and $M_2$ and virial velocity dispersions $\\sigma_{\\rm v,1}$ and $\\sigma_{\\rm v,2}$, the virial velocity dispersion of the resulting galaxy is given by \\begin{equation} \\sigma^2_{\\rm v,1+2}={M_1\\sigma^2_{\\rm v,1}+M_2\\sigma^2_{\\rm v,2}\\over M_1+M_2} \\end{equation} (by definition $\\sigma^2_{\\rm v}\\equiv 2T/M$, where $T$ is the total kinetic energy of the galaxy).  For simplicity in the formula above we considered, in the initial conditions, a negligible energy of the galaxy pair when compared to the other energies involved in the process, and no significant mass loss from the resulting system. From equation (8) it follows that $\\sigma_{\\rm v,1+2}\\leq\\,{\\rm max}(\\sigma_{\\rm v,1},\\sigma_{\\rm v,2})$, i.e., the {\\it virial} velocity dispersion cannot increase in a merging process of the kind described above. On the other hand, the FP (or the less tight Faber--Jackson relation; Faber \\& Jackson 1976) indicates that the {\\it projected, central} velocity dispersion increases with galaxy luminosity.  Then, in the dissipationless merging scenario the FP tilt can be produced only by structural and/or dynamical non--homology, since the relation between central and virial velocity dispersion depends on the structure and dynamics of the system: in particular we will discuss here the second possibility, with regard to an increase in radial orbital anisotropy with galaxy luminosity.  Note that an increase in the radial orbit amount has been claimed in the past as a natural by--product of galaxy merging, and also some observations have been interpreted in this way (Bender 1988; see also Naab, Burkert \\& Hernquist 1999, and references therein). With the aid of the explored numerical models we will try to obtain some qualitative insight in this problem: it is however clear that the results should be considered at the best qualitative indications, and that a firm answer about the role of merging in producing the FP tilt can be obtained only with N-body numerical simulations of merging galaxies (see, e.g., Capelato, de Carvalho \\& Carlberg 1995).  Summarizing, the aims of this work are the following. For what concerns the FP thickness problem, we investigate the role of radial orbit instability as a factor regulating the amount of radial anisotropy for galaxies of given luminosity, and the position, relative to the FP, of the end--products of radially unstable anisotropic models. In addition, we determine whether it is possible to reproduce the whole FP tilt with a systematic variation of radial anisotropy with luminosity (using both stable and unstable initial conditions), and what is the fate of unstable models initially forced to lie on the FP.  This paper is organized as follows.  In Section 2 we describe the basic structural and dynamical properties of the investigated models; a short description of the codes used for the numerical simulations is also given.  In Section 3 we present the results and their impact on the FP thinness problem, and in Section 4, the results are discussed focusing on the origin of the FP tilt. Finally, in Section 5 the main conclusions are summarized. ",
        "conclusions": "With the aid of numerical simulations of one and two--component galaxy models  we explored the constraints imposed by the observed thickness and tilt of the FP on the amount and distribution of radial orbital anisotropy in elliptical galaxies. The main results are summarized below.  \\begin{itemize}  \\item Remarkably, all the explored models (both one and two--component, and quite independently of the density profile) are found to be unstable when their orbital radial anisotropy is high enough to place them outside the observed FP thickness (under the assumption that their isotropic parent models lie on the FP). On the contrary all stable models lie inside the FP thickness.  \\item The end--products of one--component unstable models initially placed outside the FP fall back inside the FP: in other words, the larger is the initial displacement from the FP, the stronger is the reassessment of the model structure and dynamics. The behavior of two--component models is more varied, due to the fact that the properties of their end--products are significantly affected by the amount of mass and the distribution of the DM halo.  In particular, the end--products of models with massive (either concentrated or diffuse) DM halos remain outside the FP thickness, while models with light halos behave essentially like one--component models.  \\item Since the end--products of the unstable initial conditions are not spherically symmetric, their positions on the $(\\ku,\\kt)$ plane depend on their relative orientation with respect to the line--of--sight direction.  However, the scatter due to projection effects is in general smaller than the observed thickness of the FP, both for one and two--component models.  \\item We found that it is impossible to reproduce the whole FP tilt with radially anisotropic but stable (one and two--component) models under the assumption of constant $\\ml$ and structural homology. In other words, under these assumptions, luminous galaxies would be radially unstable well before the bright end of the FP.  \\item At variance with what happens to the end--products of unstable models initially placed outside the FP thickness (but exactly for the same reasons), the end--products of unstable models with initial conditions on the FP, fall well outside the FP itself.  \\end{itemize}    Our results lead to some speculation on the formation mechanism and evolutionary scenarios of elliptical galaxies. First, if the (unknown) formation mechanism produces galaxies with various degrees of internal radial orbital anisotropy, of which the isotropic ones constitute the ``backbone'' of the observed FP, then the most anisotropic systems would be radially unstable and would evolve into final states lying on the FP. Then no ``ad hoc'' fine tuning would be required on the amount of radial anisotropy of ellipticals at the moment of their formation. In addition, our results concerning the FP tilt could give some indications about the importance of dissipationless merging in the history of the assembly of elliptical galaxies. In fact, if Es form by hierarchical {\\it dissipationless} merging, then the very existence of the FP necessarily implies structural or dynamical non--homology of the merging end--products. The possibility that we explored in this work is that of a substantial dynamical non--homology as a function of galaxy luminosity. Our simulations show that this is not a viable possibility to reproduce the FP tilt: in this scenario the FP would be destroyed by merging. However, it should be clear that we cannot rule out the possibility that merging produces a combination of structural and dynamical effects that conspire to maintain galaxies on the FP. For this reason we are now exploring this problem, with the aid of one and two--component galaxy merging simulations.  All our results on the FP thickness and tilt seem to point toward a significant dynamical homology in real galaxies, and dynamical homology in luminous Es has been recently determined by some authors (Gerhard et al. 2001); we note also that an independent observational support for dynamical homology is given by the very evidence of the $M_{\\rm BH}$-$\\sg0$ relation, which relates a dynamical independent quantity ($M_{\\rm BH}$) with a quantity strongly dependent on anisotropy ($\\sg0$).  The fact that the scatter of the $M_{\\rm BH}$-$\\sg0$ is very small means that elliptical galaxies are basically dynamically homologous systems."
    },
    "0201/astro-ph0201128_arXiv.txt": {
        "abstract": "{We present a fit to the spectral energy distribution of \\object{OH 127.8+0.0}, a typical asymptotic giant branch star with an optically thick circumstellar dust shell.  The fit to the dust spectrum is achieved using non-spherical grains consisting of metallic iron, amorphous and crystalline silicates and water ice. Previous similar attempts have not resulted in a satisfactory fit to the observed spectral energy distributions, mainly because of an apparent lack of opacity in the 3--8 $\\mu$m region of the spectrum. Non-spherical metallic iron grains provide an identification for the missing source of opacity in the near-infrared.  Using the derived dust composition, we have calculated spectra for a range of mass-loss rates in order to perform a consistency check by comparison with other evolved stars. The $L-[12$ $\\mu$m$]$ colours of these models correctly predict the mass-loss rate of a sample of AGB stars, strengthening our conclusion that the metallic iron grains dominate the near-infrared flux. We discuss a formation mechanism for non-spherical metallic iron grains. ",
        "introduction": "\\label{sec:introduction}  Oxygen-rich Asymptotic Giant Branch (AGB) stars show an infrared excess on their spectral energy distribution (SED), which arises from thermal emission of dust located in a circumstellar shell. It is well established that the dust in this shell mainly consists of silicates, deduced from the clear detection of the 10 and 18 $\\mu$m features. These bands are due to the Si-O bond stretching and Si-O-Si bond bending vibrations, respectively.  It was noticed already some 25 years ago \\citep{JM_76_dust,B_77_dustshells} that the opacities of various types of silicates measured in the laboratory, were not high enough to explain the shape of the spectrum in the near-infrared (NIR, $3 < \\lambda < 8$ $\\mu$m). \\citet{JM_76_dust} investigated the effect of pure silicates with graphite inclusions and of meteoritic rock on the shape of the SED. These materials were taken to represent silicates with various metallic inclusions, and referred to as \\emph{dirty silicates}. Radiative transfer calculations using the opacity of dirty silicates were in good agreement with the observations.  \\citet{DL_84_optprop} used laboratory measurements of limited resolution and wavelength coverage \\citep{HS_73_optpropsolids} in which the optical constants are partly modified and expanded in wavelength regime to match the astronomical observations of interstellar and circumstellar silicates \\citep{JM_76_dust,RMC_83_mucep} These optical constants are usually referred to as \\emph{astronomical silicate}.  The exact chemical composition of this astronomical silicate is not known, nevertheless it is widely used in theoretical studies of the radiative transfer in circumstellar dust shells \\citep[e.g.][]{B_87_dustshells,ST_89_theory,JT_92_SED,LL_93_SED}. Recently there have been efforts to derive improved optical constants for astronomical silicates, both for lines-of-sight in the interstellar medium (ISM) and toward oxygen-rich evolved stars exclusively \\citep{OHM_92_cosmicsilicates,DP_95_cssil,S_99_optprop}. These optical constants still do not provide the answer to the composition and the nature of circumstellar silicates.  The study of circumstellar dust shells in the pre-ISO era was based on low resolution spectroscopy, but with the launch of the Infrared Space Observatory (ISO) \\citep{KSA_96_ISO}, a large wavelength region, from 2--200 $\\mu$m, became available for intermediate resolution spectroscopic observations ($\\lambda / \\Delta \\lambda \\gtrsim 400$). The ISO spectra of oxygen-rich evolved stars turned out to be extremely rich in solid state features. Hence, it became meaningful to attempt spectral fits using properly characterized cosmic dust analogs, rather than astronomical silicates. \\citet{DDW_00_OHIR} aim to do this for OH/IR stars (i.e.~AGB stars with an optically thick dust shell), but their work is not complete and parts of the spectrum are not fitted to a satisfactory level.  \\citet{HMD_01_VYCMa} have studied the optically thin dust shell surrounding red supergiant \\object{VY CMa} and are able to fit the SED, albeit after including a large overabundance of metallic iron in the dust composition.  In this study, we aim to determine the dust composition with full radiative transfer calculations using exclusively optical constants of well-defined materials measured in the laboratory, examining all astronomically relevant materials, and taking abundance constraints into account. Revealing the exact dust composition is an important step toward understanding dust formation and processing in the outflows of oxygen-rich evolved stars.  The paper is organized as follows: In Sect.~\\ref{sec:near-infr-probl} we discuss the spectral energy distribution arising from full radiative transfer calculations using amorphous olivine as the only dust component, and compare it to the observations. Sect.~\\ref{sec:improvedfit} discusses the dust composition and grain properties required to improve the fit to the observations. A spectral fit to \\object{OH 127.8+0.0} is presented in Sect.~\\ref{sec:oh127}, supported by a consistency check on other AGB stars in Sect.~\\ref{sec:lmin12}.  Section~\\ref{sec:disc-concl} contains a discussion of the results. ",
        "conclusions": "\\label{sec:disc-concl}  The results presented in this work provide the first sound identification of metallic iron in the circumstellar dust shell around AGB stars, without violating abundance constraints. \\citet{DDW_00_OHIR} have tried to include metallic iron in their model calculations of evolved stars, but because they did not consider non-spherical grains, metallic iron could not be abundant enough to account for the missing NIR opacity. \\citet{HMD_01_VYCMa} have concluded that metallic iron is necessary to fit the spectrum of \\object{VY CMa}, although all Fe-atoms should be incorporated in metallic iron, implying that the amorphous silicate is completely Mg-rich, and even then a large overabundance is required.  From meteoritic and theoretical studies it is known that metallic iron can condense in a cooling gas of solar composition \\citep[e.g.][]{GL_74_chemicalhistory,LN_79_iron,KH_88_ironbearing}. In that case, metallic iron forms almost simultaneously with silicates, since its condensation temperature is only 50 -- 100 K below that of silicates.  In an oxygen-rich chemistry, metallic iron is stable above 700 K. Below that temperature it will react to form FeS and/or FeO \\citep{J_90_iron}.  An alternative is that it will remain metallic because it is protected from the oxidizing environment by inclusion in a grain. This can be achieved through the formation of \\emph{iron islands} on the surface of silicate grains \\citep{GS_99_condensation}. During the formation process of silicates, Fe-atoms in the lattice will be replaced by the thermodynamically more favourable Mg. The Fe-atoms will migrate to the surface of the grain to form \\emph{iron islands} and stimulate the condensation of additional Fe-atoms, to continue the growth of the islands.  \\citet{GS_99_condensation} suggest that the islands will eventually be covered by younger silicate layers, thus creating silicate grains with platelet shaped, i.e.~non-spherical, metallic iron inclusions. These inclusions greatly increase the opacity of the grains in the NIR region.  The model calculations presented in this work use separate grain populations, rather than metallic iron inclusions in silicate grains, but still provide an idea of the mass fraction contained in the iron inclusions.  This result, together with the inclusion of crystalline water ice and the use of CDE for all grain shapes provide an important step in disentangling the dust composition of the circumstellar shell of AGB stars.  No exotic assumptions were required to match the spectrum. A spherical outflow of material at constant velocity suffices to produce a good fit to the SED \\emph{if} chemical and shape properties of the grains are treated in sufficient detail.  The presence and formation of metallic iron in astrophysical environments has been subject to many studies in the past \\citep[][and references herein]{J_90_iron}.  Studies of the interstellar extinction toward the \\object{galactic centre} (GC) revealed a discrepancy in the 3--8 $\\mu$m region with the standard dirty silicate opacity \\citep{L_99_gc,LFG_96_gc}. Of course, this is an ISM line-of-sight, implying that amorphous carbon is an important dust component. On the other hand, the opacity contribution of amorphous carbon is already included in the optical constants of the dirty silicate, i.e.~a graphite-silicate mixture. Thus also in this case, non-spherical metallic iron grains might explain the missing opacity in the NIR region.  In addition, it is known that \\emph{superparamagnetic} inclusions, such as metallic iron, in elongated silicate grains can cause these grains to align when a magnetic field is present \\citep{M_86_alignment}.  The alignment of silicate grains explains the polarization of starlight passing through the interstellar medium, a suggestion first made by \\citet{ST_51_polarization}.  This theory appears to be supported by studies of interplanetary dust grains, which are collected by aircraft in the upper atmosphere of the Earth.  Some of the collected particles are believed to be very pristine and probably even from interstellar origin, rather than being reprocessed during the formation of the solar system. These pristine particles consist of an amorphous matrix material with metal inclusions, and are referred to as GEMs (glasses with embedded metal and sulphides).  Metallic iron inclusions are quite common in these supposedly interstellar grains \\citep{B_94_anomalousIDP,M_95_GEMs}. Whether these grains really originate from the outflows of evolved stars remains subject to speculation at this moment."
    },
    "0201/astro-ph0201364_arXiv.txt": {
        "abstract": "{We compare the 3D clustering of old passively-evolving  and dusty star-forming $z\\sim1$ EROs from the K20 survey. With detailed simulations of clustering, the comoving correlation length of dusty star-forming EROs is constrained to be less than $r_0 \\sim 2.5$ \\h1 Mpc. In contrast, the old EROs are much more positively correlated, with $5.5 \\simlt r_0/($\\h1 Mpc$) \\simlt 16$, consistent with previous claims for $z\\sim1$ field early-type galaxies based on analyses of ERO angular clustering. The low level of clustering of dusty star-forming EROs does not support these to be major mergers building up an elliptical galaxy, or typical counterparts of SCUBA sources, but it is instead consistent with the weak clustering of high redshift blue galaxies and of luminous local IRAS galaxies. Current hierarchical merging models can explain the large $r_0$ for $z\\sim1$ field early-type galaxies, but fail in matching their high number density and overall old ages. ",
        "introduction": "Extremely red objects ($R-K>5$, EROs hereafter) are providing increasingly stringent constraints on our understanding of the formation of galaxies in general, via their spectral evolution and clustering properties. The very red colors of EROs are well known to be both consistent with old passively evolving distant ($z>0.8$) elliptical galaxies (e.g. Cohen et al. 1999; Spinrad et al. 1997) or dust-reddened starburst galaxies (e.g. Cimatti et al. 1998; Smail et al. 1999). Purely passive evolution of the present day population of elliptical galaxies is consistent with the measured surface density of faint EROs with $K\\sim17$--22, while current renditions of the semianalytical hierarchical merging models fail to reproduce the surface density of EROs by a large factor (Daddi et al. 2000a; Smith et al. 2001; Firth et al. 2001).  Recently, we have completed a relatively large deep survey of very red galaxies covering 700 arcmin$^2$ (Daddi et al. 2000b, D00 hereafter), concluding that EROs are strongly clustered in projection, by an order of magnitude more than all galaxies at the same limits of $K\\leq 18$--19.2. With careful attention to the measurement uncertainty inherent in narrow field data, Daddi et al. (2001, D01 hereafter) showed that the angular clustering of EROs implies a spatial correlation length of $r_0=12\\pm3$ \\h1 comoving Mpc, consistent with the assumption that the ERO population is dominated by elliptical galaxies. This large clustering amplitude is not at odds with recent hierarchical merging models, which require that the most massive galaxies are clustered more strongly than the general galaxy population at high $z$ (e.g. Mo \\& White 1996). Our results on the angular and spatial clustering of EROs have been substantially confirmed by the Las Campanas Redshift Survey data (McCarthy et al. 2001; Firth et al. 2001; Moustakas \\& Somerville 2001).  {In our recent large K20 redshift survey of a flux limited sample of $\\sim500$ galaxies with $K\\leq20$} (Cimatti et al. 2002, C02 hereafter), we obtained redshifts for a sub-sample of 35 EROs. For red objects with $R-K>5$ and $K\\leq19.2$, about 1/3 were identified as old systems (consistent with being passively evolving elliptical galaxies), 1/3 were found to be dusty starburst galaxies and 1/3 remain unidentified. While the derived fraction of early-type galaxies, $50\\pm20$\\%, is consistent with previous estimates based on morphology (Moriondo et al. 2000; Stiavelli \\& Treu 2000), C02 showed that the dusty star-forming (SF) objects do contribute significantly to the ERO population at faint magnitudes, thus complicating the interpretation of both ERO surface density and clustering, as measured in earlier analyses.  In particular, given the strong interest in the clustering amplitude of early-type galaxies, it is important to estimate separately the clustering properties of the old and of the dusty-SF EROs, hence their relative contribution to the clustering of the whole ERO population. This is attempted in this letter, where we adopt a cosmology with $\\Omega_\\Lambda = 0.7$, $\\Omega_m = 0.3$ and $H_0 = 100h$ km/s/Mpc. ",
        "conclusions": "\\subsection{The clustering of dusty-SF EROs}  The clustering of dusty-SF EROs is small, and maybe consistent with the values of $1\\simlt r_0/($\\h1 Mpc$) \\simlt 2.5$ measured for star-forming galaxies at $z\\sim1$ (Le Fevre et al. 1996, Carlberg et al. 1997, Hogg et al. 2000). This would suggest the former to be a subclass of the latter, but with stronger dust extinction. Locally, dusty galaxies detected by IRAS are also known to have a relatively weak clustering (e.g. Saunders et al. 1992).  The low level of clustering seems also to be at odds with the idea that the dusty-SF EROs are in a starburst phase following a major merger event, eventually expected to produce an elliptical galaxy, as in this case one would expect to find a correlation length somewhat lower than, but similar, to that of the ellipticals at the same redshift.  SCUBA sub-mm selected sources are also thought to be dusty objects at high redshift detected by virtue of the emission from dust warmed by star-formation or AGN activity. This population is expected (Magliocchetti et al. 2001) and tentatively observed (Scott et al. 2001) to show strong angular clustering at the level of $A(1^o)\\sim0.01$ (see also Ivison et al. 2000 and Chapman et al. 2001). Therefore our result suggest the dusty-SF EROs are a different population with respect to SCUBA galaxies, with small overlap, in agreement with the latter being typically fainter and more distant ($K>20$, median redshift $z\\sim2.5$--3, {Smail et al. 2000}, ; see also C02, Mohan et al. 2001, {and Dannerbauer et al. 2002 for MAMBO sources}).   Finally, we note that the class of dusty-SF EROs could be internally inhomogeneous: some of them maybe spirals with moderate extinction (e.g. van Dokkum \\& Stanford 2001) and there could also be a mixture of dust-enshrouded AGNs and starburst galaxies (C02).  \\subsection{Field $z\\sim1$ early-type galaxies}  {Observations of samples of faint ERO galaxies have led to three key conclusions regarding bright early-type galaxies with $L\\simgt L_*$, at $z\\sim1$. Firstly, their space density is consistent with that of local luminous early-type galaxies, when account is made of minimal pure luminosity evolution (PLE) (C02). Secondly, spectroscopy implies age $\\simgt3$ Gyr for their stellar populations (assuming solar metalicity, C02); and thirdly, a comoving correlation length $r_0\\simgt 12$ \\h1 Mpc (this paper and D01) has been measured comparable with the local value for luminous early-type galaxies.}  A large correlation length, $r_0 \\simgt 10$ \\h1 Mpc, is anticipated theoretically for the hierarchical merging paradigm for which a rapidly increasing bias is predicted for massive galaxies by $z\\sim1$ (e.g. Mo \\& White 1996, Moscardini et al. 1998).  {Such a large correlation length is not expected for the PLE (galaxy conservation) scenario (D01).} However, also current semi-analytical renditions of the hierarchical models seem to be at odds with the observed results. For example, the Cole et al. (2000) model predicts a comoving density (Fig. 1 of Benson et al. 2001) of {\\em all} the $z\\sim1$ galaxies with $10^{11}M_\\odot$ (consistent with our $K\\leq19.2$ selection) which is a full order of magnitude below the density of just the old EROs observed by C02. Similarly, the Kauffmann et al. (1999) model \\footnote{http://www.mpa-garching.mpg.de/GIF/} predicts a comoving density of $z\\sim1$ EROs ($R-K\\geq5$, $K\\leq19.2$) that is 3(6) times lower than observed by C02 for old(all) EROs.  In addition, in these models $z\\sim1$ galaxies qualified as field early-types appear to have experienced recent star-formation, while the present sample of old EROs is dominated by an old stellar population.  We conclude that to our knowledge no semianalytical rendition of the hierarchical merging models can yet account for all the 3 key observed properties of $z\\sim1$ field early type galaxies described above."
    },
    "0201/astro-ph0201478_arXiv.txt": {
        "abstract": "{ We combine the radio continuum images from the NRAO VLA Sky Survey with the CO-line observations from the extragalactic CO survey of the Five College Radio Astronomy Observatory to study the relationship between molecular gas and the star formation rate within the disks of 180 spiral galaxies at 45\\arcsec~resolution. We find a tight correlation between these quantities. On average, the ratio between the radio continuum and the CO emission is constant, within a factor of 3, both inside the same galaxy and from galaxy to galaxy. The mean star formation efficiency deduced from the radio continuum corresponds to convert 3.5\\% of the available molecular gas into stars on a time scale of $10^{8}$ yr and depends weakly on general galaxy properties, such as Hubble type or nuclear activity. A comparison is made with another similar analysis performed using the \\ha luminosity as star formation indicator. The overall agreement we find between the two studies reinforces the use of the radio luminosity as star formation rate indicator not only on global but also on local scales. ",
        "introduction": "One of the major goals of the studies of external galaxies is understanding the relationship between the star formation rate (SFR) and the physical condition in the interstellar medium. Several indicators have been suggested to estimate the SFR (of massive stars) in galaxies. These include the U-band magnitude, the strength of Balmer lines emission, the far-infrared (FIR) emission and the radio luminosity; the rates inferred from the different indicators span almost four orders of magnitude going from $10^{-2}$ to $10^{2}$ \\msyr. Cram et al. (1998) checked  the consistency between the SFR deduced from the U-band, H$_{\\alpha}$, FIR and radio luminosity using a sample of 700 local galaxies. They noted that there are systematic differences between these various indicators. In particular they suggested that the H$_{\\alpha}$ luminosity may underestimate the SFR by  approximately an order of magnitude when the SFR is $\\geq 20$ \\msyr. They concluded that the radio continuum luminosity at decimeter wavelengths of a star forming galaxy provides a better way to estimate the current rate of star formation. The radio continuum emission at 1.4 GHz from a star-forming galaxy is mainly synchrotron radiation produced from relativistic electrons accelerated by supernovae explosions (Lequeux 1971). Indeed the radio continuum luminosity appears to be directly proportional to the rate of formation of supernovae (Condon 1992). This view is reinforced by the tight correlation existing between the radio luminosity and the FIR for spiral galaxies (see e.g. Condon 1992). Since the radio continuum at 1.4 GHz does not suffer significant extinction, the radio luminosity constitutes a very useful tool to determine the current SFR in a spiral galaxy.  Since the discovery that stars form in molecular clouds, it is essential to determine, not only the rate, but also the efficiency of conversion of the interstellar gas in stars; i.e. the star formation efficiency (SFE). The SFE measures the formation rate of young stars per unit of mass of gas available to form those stars. Determining the SFE is important to distinguish a situation in which a high SFR indicates a higher efficiency in converting gas in stars rather than a higher gas quantity.  The CO molecule luminosity  and the virial mass of giant molecular clouds correlate very well in our Galaxy and in other nearby spirals (Young \\& Scoville 1991 and references therein). The comparison of different SFR tracers with the mass of molecular clouds provides indeed an important tool to investigate the behaviour of the SFE within and among galaxies. Many studies have been concerned with the behaviour of the star formation process on global scales, averaged over the entire star-forming disk. These works showed that the disk-averaged star formation process is well described by a Schmidt (1959) law of the type $\\Sigma_{\\rm SFR} \\propto  \\Sigma_{\\rm gas}^{N}$, where $\\Sigma_{\\rm SFR}$ and $\\Sigma_{\\rm gas}$ are the observable surface density of SFR and total (atomic + molecular) gas density, respectively, and the exponent $N$ typically ranges from 1.3 to 1.5 (Kennicutt 1998).  An interesting development of these global studies, the investigation of the behaviour of the SFE {\\it within} the disks of the individual galaxies, provides much physical insight into the star formation process itself. The extragalactic CO survey of the Five College Radio Astronomy Observatory (Young et al. 1995, hereafter FCRAO CO Survey) provided a uniform database of CO data for 300 galaxies at a resolution of 45\\arcsec, opening the possibility to extend the study of the Schmidt relationship of the SFR versus the H$_{2}$ density over the same physical regions well inside the galaxy disks. Since the star formation process involves the molecular gas directly, some authors recognized that the determination of the Schmidt law assumes a clear physical meaning if restricted to this gas component. Moreover, in the considered regions the molecular gas is dominant over the atomic one and, contrary to this latter, its azimuthally averaged distribution follows closely the radial profiles of the main SFR indicators (Tacconi and Young 1986, Young \\& Scoville 1991). Rownd \\& Young (1999; hereafter RY99) conducted an H$_{\\alpha}$ imaging of 121 of these galaxies, determining the local relationship between the SFR and the molecular gas. They found a correlation between these two quantities and concluded that for face-on spiral, in general, there are no strong SFE gradients across the star-forming disks. The majority of large SFE variations they found are seen between adjacent disk points, reflecting regional differences in the SFE, and any radial gradients are at most a secondary effect. In contrast, they pointed out that consistent radial variations (up to an order of magnitude or more) of the SFE exist within many highly inclined galaxy disks. They attributed the decreasing SFE towards the centers of these galaxies to a large amount of dust extinction on the \\ha luminosity.  Adler et al. (1991) found a correlation between the radio continuum flux density at 20 cm and the CO line emission on global scales for a sample of 31 spiral galaxies. They also studied the relationship of these two quantities within the disks of 8 nearby well resolved spiral galaxies, finding that their ratio is constant both inside the same galaxy and from galaxy to galaxy.\\\\  The work we present here is complementary to the analysis of RY99 and extends that of Adler et al. (1991). We combined the radio continuum images at 1.4 GHz from the NRAO VLA Sky Survey (NVSS, Condon et al. 1998) with the FCRAO CO survey to study the relationship between the radio continuum and the molecular gas point-to-point within the disks of 180 star-forming spiral galaxies. It is important to stress that we are comparing two homogeneous data set with the same angular resolution of 45\\arcsec.\\\\  The paper is organized as follows: in Sect.~2 and Sect.~3 we present the sample used and we describe the data analysis, respectively. In Sect.~4 we present the results of the statistical analysis and in Sect.~5 we discuss the results obtained.\\\\  We use a Hubble constant H$_0$=50~km~s$^{-1}$Mpc$^{-1}$ throughout the paper. ",
        "conclusions": "The most striking result of our study is the overall consistency of the SFE obtained from the 20~cm radio continuum as SFR indicator instead of the \\ha~emission (RY99). Basically we found all the essential features observed by RY99: the constancy of the SFE both within and among galaxies; the weak dependence of the SFE on morphological type and linear size of the observing beam; the extent and slope of the composite Schmidt law. Nevertheless, some important differences arise. In this section we present in details the differences between the two works, discussing their nature and implications.  \\begin{figure}[t] \\begin{center} \\includegraphics[angle=0, width=8.5cm]{fig10.ps} \\end{center} \\caption[]{ Comparison of SFRs surface densities deduced from 1.4 GHz luminosity (horizontal axis) and \\ha emission (vertical axis) for 102 galaxies of our sample in common with RY99. The various panels show different values of galaxy inclinations $i$, as indicated. There are 453 detections, 65 of which are upper limits (at 2$\\sigma$). The reference dashed lines indicate 1:1 relations. The subsample of face-on galaxies ($i \\leq 40\\degr$) shows the best correlation, with a dispersion of a factor of 2.3. SFRs surface densities deduced by the \\ha for highly inclined galaxies ($i \\geq 70\\degr$) are systematically underestimated for $\\Sigma_{\\rm SFR} {\\rm 1.4~GHz} \\geq 3\\times 10^{-3} ~{\\rm M_{\\odot}~yr ^{-1}kpc^{-2}}$.} \\label{fig10} \\end{figure}  \\subsection{Star formation rate surface densities} Cram et al. (1998) checked the consistency between the SFR deduced from the radio luminosity with the rates predicted by other indicators, in particular by the \\ha luminosity. They concluded that the rates deduced from the radio continuum and the \\ha luminosities, although in broad agreement, are affected by two systematic deviations: respectively at low ($\\leq 0.1 ~{\\rm M_{\\odot}yr^{-1}}$) and high ($\\geq 20 ~{\\rm M_{\\odot}yr^{-1}}$) star formation rates, the \\ha luminosity overestimate and underestimate the SFR deduced by the radio luminosity. The deviation at low rates can be attributed in part to problems related to the zero-point \\ha luminosity calibration. They suggest that the deviation at high SFR could be attributed to a larger amount of extinction by dust in those objects undergoing strong star formation or to particular IMFs that favour low mass supernovae progenitors rather than high mass stars responsible for the \\ha.  By comparing our data set with that of RY99 we have the possibility to extend the consistency check between the surface star formation rate densities deduced from the radio continuum and  \\ha emission over the same regions of galaxy disks.  \\begin{figure}[b] \\begin{center} \\includegraphics[angle=0, width=9cm]{fig11.ps} \\end{center} \\caption[]{ Ratio of SFRs surface densities deduced from \\ha and 1.4 GHz luminosity for  highly inclined (NGC1055 and NGC3079) and starburst galaxies (NGC2146, NGC3034). These examples dramatically illustrate that the SFRs surface density inferred from the \\ha emission towards the nuclei of these objects are underestimated by more than one order of magnitude.} \\label{fig11} \\end{figure} Using Eq.~(2) of Cram et al. (1998), we calculated the SFR surface density from the \\ha brightness reported by RY99 through the formula: \\begin{equation} \\Sigma_{\\rm SFR}\\left({\\rm \\frac{M_{\\odot}}{yr~kpc^{2}}}\\right) = 2.6\\times10^{-2}~\\mu({\\rm H_{\\alpha}}) \\label{hasfr} \\end{equation} where the (de-projected) \\ha surface brightness, $\\mu({\\rm H_{\\alpha}})$, is measured in ${\\rm L_{\\odot}~pc^{-2}}$. Eq.~(2) of Cram et al. (1998) accounts for a correction of 1.1 magnitude in the luminosity to compensate for the mean extinction in \\ha (Kennicutt 1983a). Fig.~10 shows the direct comparison between the radio and \\ha SFRs inferred respectively from Eq.~(\\ref{sfrb}) and Eq.~(\\ref{hasfr}) for 453 position within 102 galaxies of the sample. All galaxies are shown in the top panel of Fig.~10. In general there is a broad agreement between the two SFR indicators. Considering the distribution of the logarithm values of the  ratio $R=\\Sigma_{\\rm SFR(H_{\\alpha})}/\\Sigma_{\\rm SFR(1.4 GHz)}$, we have that $R$ have a mean of 0.83 with a standard deviation of 2.8, i.e. the \\ha is underestimating the surface SFR as compared to continuum radio emission. Sorting the sample by galaxy inclination indicates that this effect is strongly correlated with galaxy orientation. For the face-on subsample ($i < 40\\degr$) $R$ has a mean of 1.1 and a dispersion of a factor 2.3. At intermediate inclinations ($40\\degr \\leq i < 70\\degr$) mean and dispersion of $R$ are 0.8 and a factor of 3, respectively. Finally,  for highly inclined galaxies ($i \\geq 70\\degr$)  mean and dispersion of $R$ are  0.7 and a factor of 2.9, respectively. In particular, SFR surface densities deduced by the \\ha for highly inclined galaxies are systematically underestimated for $\\Sigma_{\\rm SFR(1.4~GHz)}$ greater than  $\\sim 3\\times 10^{-3} ~{\\rm M_{\\odot}~yr^{-1}kpc^{-2}}$. Fig.~11 shows two edge-on spirals, NGC3079 and NGC1055 (see Figs. 2 and 3) and two well known starburst galaxies, NGC2146 and M82, in which  SFRs deduced by the \\ha emission are underestimated by more than a factor of about 10. As a consequence, the SFE variations deduced for these objects by RY99, on the basis of \\ha observations, can be attributed to extinction, as suspected by these authors.  These results have two important implications: i) the close correlation observed between $\\Sigma_{\\rm SFR(1.4~GHz)}$ and $\\Sigma_{\\rm SFR(H_{\\alpha})}$ for face-on galaxies reinforces the use of the radio luminosity as SFR indicator not only on global but also on local scales; ii) extinction could significantly affect estimates based on \\ha emission for high SFRs in edge-on galaxies.  Although the star formation rates deduced from the \\ha emission are systematically underestimated compared to those deduced from the radio continuum, the mean SFE reported by RY99 for their entire sample (121 galaxies) is 4.3\\%, i.e. higher than the mean SFE deduced from the radio continuum for our entire sample of 180 galaxies. However, considering the 103 galaxies we have in common with RY99, the mean SFE deduced from the \\ha emission is about 3.1\\% in good agreement with our value.  \\begin{figure}[h] \\begin{center} \\includegraphics[angle=0, width=9cm]{fig12.ps} \\end{center} \\caption[]{ Ratio of SFRs surface densities deduced from 1.4 GHz luminosity and \\ha emission as functions of the linear scale of the beam (top panel) and distance from galaxy center (bottom panel). Highly inclined galaxies ($i> 70\\degr$) are represented with open dots.} \\label{fig12} \\end{figure} \\subsection{Non-thermal radio continuum scale-length} Discrete supernova remnants (SNR) themselves account only for $<10\\%$ of the radio luminosity of normal galaxies (Ilovaisky \\& Lequeux 1972). Most of cosmic rays escape from their parent SNRs and diffuse along the galaxy disk during their characteristic lifetimes which are much larger than typical ages of SNRs ($\\sim 10^{5}$ yr). Therefore one expects the non-thermal radio continuum emission to be smoothed over a larger region around the parent stellar population with respect to the \\ha emission. On scale smaller than the cosmic rays characteristic diffusion scale, $D_{\\rm CR}$, the non-thermal radio continuum cannot provide a reliable estimate of the local SFR. On the other hand, averaging over regions larger than the typical spiral arm width leads to a dilution of the \\ha emission. Fig.~12 (top panel) shows the ratio between the SFR surface densities  deduced  from \\ha and from radio luminosity versus linear size of the observing beam. One notes that, apart for edge-on galaxies (where the \\ha is affected by extinction), the ratio $\\Sigma_{\\rm SFR(1.4~GHz)}$ to  $\\Sigma_{\\rm SFR(H_{\\alpha})}$ tends to increase with the beam size. In particular, the ratio is smaller or greater than one for beam sizes respectively below and above $\\sim 5$ kpc. This is consistent with a scenario in which the characteristic linear scale of both cosmic rays diffusion scale and arm size is about 3 kpc. Higher resolution radio images are required to determine the lowest linear scale at which the radio continuum  correlates with the \\ha emission and thus providing an estimate of $D_{\\rm CR}$.  We found also that the ratio  $\\Sigma_{\\rm SFR(1.4~GHz)}$ to $\\Sigma_{\\rm SFR(H_{\\alpha})}$ on average does not depend on the distance from the galaxy centers (Fig.~12, bottom panel).  \\subsection{The nature of Schmidt law scatter} The dispersion (standard deviation) of the composite radio Schmidt law shown in Fig.~9 is a factor of 2.8, which is smaller than the corresponding dispersion of a factor 4 of the \\ha Schmidt law (RY99). This is probably because the radio continuum does not suffer the effects of dust extinction. The noisiness of the correlation could reflect real variations of the mean Schmidt law, indicating that it should be regarded at best as a first  order approximation of the star formation process, or it could be due to the many assumptions involved in the calculation of gas and SFR densities.  Variations of the $X_{\\rm CO}$ factor from galaxy to galaxy can be advocated to explain a part of the correlation scatter. These could introduce uncertainties up to a factor of 2 in the gas density scale (Kennicutt 1998). Another possibility is that the extrapolation of the proportionality between CO luminosity and virial mass of giant molecular clouds observed in our own Galaxy and in nearby galaxies, which is the basis of molecular mass determinations in this and similar works, does not hold for all spiral galaxies. \\begin{figure}[t] \\begin{center} \\includegraphics[angle=0, width=8.75cm]{fig13.ps} \\end{center} \\caption[]{ Radio Schmidt law for Sa, Sbc and Sc (top panel) and S0-Sab, Scd, Sd-Sm and merging-irregular (bottom panel); limits are not shown. Dashed and dotted reference lines indicate the mean and the dispersion of the two sub-samples, respectively. The scatter of the Schmidt law is considerably reduced excluding extreme morphological types. The scatter of the correlations shown in top and bottom panels is a factor of 2.4 and 3.8, respectively. In top panel, the deviation from the correlation of the four points belonging to NGC1068 is due to the AGN-related radio emission of this Seyfert galaxy (see Sect.~4.2).} \\label{fig13} \\end{figure}  However, by comparing SFR surface densities deduced from \\ha and from radio continuum luminosity we showed that, even in the face-on subsample, the $\\Sigma_{\\rm SFR(H_{\\alpha})}-\\Sigma_{\\rm SFR(1.4~GHz)}$ relation itself is affected by a scatter of at least a factor of 2. Hence, the uncertainties of the SFR indicators alone can account for a consistent fraction of the Schmidt law scatter.  Fig.~13 shows the Schmidt law separately for Sb, Sbc and Sc and  S0-Sab, Scd and Sd-Sm galaxies (only detections with signal-to-noise ratio greater than 2$\\sigma$ are considered). The two subsamples have the same mean SFE of 3.5\\% but very different dispersions of a factor of 2.4 for the former and 3.8 for the latter, i.e. the scatter of the Schmidt law is considerably reduced excluding extreme morphological types. RY99 also found that the \\ha Schmidt law is considerably tightened with the exclusion of the irregular galaxies and mergers. These SFE variations around the mean Schmidt law can be attributed to the particular physical conditions and/or environments experienced by these objects (see e.g. starburst galaxies), but they could also be induced by observational effects further amplified by poor statistic. In fact, Scd galaxies which are characterized by the lower mean SFE in our sample (see also Fig.~5) behave consistently  with other morphological types excluding IC342 for which the NVSS misses flux density from the extended structure, see Sec.~2.  \\subsection{Molecular gas consumption timescale} The star formation efficiency can be expressed in terms of gas consumption (or cycling) timescale. The molecular gas consumption timescale, $\\tau_{\\rm H_{2}}$, indicates the time needed to convert all the molecular gas in stars given a constant SFE. Expressed in Gyr, $\\tau_{\\rm H_{2}}$ is given by \\begin{equation} \\tau_{\\rm H_{2}}({\\rm Gyr})=\\frac{10}{\\rm SFE_{\\tau 8} (\\%)} \\end{equation} The mean star formation efficiency found in this work for spiral galaxies yields a molecular gas consumption timescale of about $\\tau_{\\rm H_{2}} \\simeq 2.8$ Gyr. Given the dispersion of the SFE in the sample, most galaxies have gas cycling timescale between about 1 and 9 Gyr. At the lower end of  the distribution we found galaxies with a \\sfe$\\sim 1\\%$, such as NGC4535 (see Fig.~2). The corresponding molecular gas consumption timescale is almost comparable with the Hubble time. The shortest gas cycling timescale corresponds to galaxies with exceptional SFE like NGC3310 (see Fig.~3). Star formation rates in these galaxies are so high, compared to their gas supplies, that the gas cycling timescale is  $\\tau_{\\rm H_{2}} \\le 0.1$ Gyr.  It is interesting to investigate the behaviour of the so-called starburst galaxies with respect to the gas cycling timescale. In literature, galaxies have been classified as starbursts according to different criteria. Heckman et al. (1998) define a galaxy as starburst when it is hosting a star-forming event that dominates its bolometric luminosity, i.e. on the basis of the magnitude of the SFR. Alternatively, Shu (1987) and Young (1987) proposed a classification based on the efficiency of the star formation. In this latter definition a galaxy with a high SFR is not defined as starburst if the mass of gas available is enough to sustain the star formation rate and vice versa. Following RY99 we selected the galaxies hosting a region in which the SFE is enhanced by a factor of three compared to the mean of the sample; for these starburst regions $\\tau_{\\rm H_{2}} \\le 1$ Gyr. Fig.~14 shows the usual $\\Sigma_{{\\rm SFR}} - \\Sigma_{{\\rm H_{2}}}$ plane for all the detection in the sample along with three reference lines indicating the gas cycling timescales $\\tau_{\\rm H_{2}}=0.1$, 1 and 10 Gyr. \\begin{figure}[t] \\begin{center} \\includegraphics[angle=0, width=9cm]{fig14.ps} \\end{center} \\caption[]{ $\\Sigma_{{\\rm SFR}} - \\Sigma_{{\\rm H_{2}}}$ plane for all the detection in the sample (limits are not shown) along with three reference lines corresponding to different molecular gas consumption timescales, as indicated. Pointings of known starburst galaxies are represented by open dots. } \\label{fig14} \\end{figure} Most galaxies known in literature as ``starbursts'' have consumption timescales comparable with those of normal spiral galaxies. In some cases, e.g. M82, starburst galaxies host both regions characterized by SFE lower and higher than the mean of the sample. On the other hand, galaxies with normal or even low SFR surface densities, such as the highly inclined galaxy NGC4631, should be regarded as starburst according to our selection criterion based on the SFE. Finally, there are some galaxies for which the SFE is so high that the gas cycling timescale is $\\tau_{\\rm H_{2}} \\le 0.1$ Gyr, e.g. the starburst NGC3310. NGC3310 is one of the best examples of a local UV-bright starburst (see Conselice et al. 2000 and reference therein). The exceptional high SFE of this galaxy indicates that its interstellar medium is truly affected by an extraordinary star formation event.  \\subsection{Conclusions} The results of this work are the following:  1.~There is a tight correlation between the 20 cm non-thermal radio continuum and the CO line intensity in a representative sample of 180 spiral galaxies. The correlation holds within and among the galaxies.  2.~The mean star formation efficiency, i.e. the ratio of the radio SFR to the molecular gas densities,  for our sample is $0.035\\times 10^{-8}$ yr$^{-1}$ with a dispersion of a factor of 3. This corresponds to convert 3.5\\% of the available gas into stars on a time of $10^{8}$ yr.  3.~The comparison of SFRs surface densities deduced from 1.4 GHz luminosity and from the \\ha emission for 102 galaxies, reveals that $\\Sigma_{\\rm SFR(1.4~GHz)}$ and $\\Sigma_{\\rm SFR(H_{\\alpha})}$ are closely correlated for face-on galaxies ($i \\leq 40\\degr$), reinforcing the use of the radio radio luminosity as SFR indicator not only on global but also on local scales. SFRs surface densities deduced by the \\ha luminosity for highly inclined galaxies ($i \\geq 70\\degr$) are systematically underestimated for $\\Sigma_{\\rm SFR(1.4~GHz)} \\geq 3\\times 10^{-3} ~{\\rm M_{\\odot}~yr ^{-1}kpc^{-2}}$.  4.~The  star formation efficiency varies weakly (less than 25\\%) with the Hubble morphological type.  5.~The variation of the SFE within individual galaxy disks is less than a factor of 3. The largest variations are found in starburst galaxies.  6.~The SFE is found to be approximately constant as a function of distance from the galaxy centers.  7.~The composite radio Schmidt law, star formation versus molecular gas content,  extends for more than 3 order of magnitude with an exponent of 1.3.  8.~Most galaxies known in literature as ``starbursts'' have consumption timescales comparable with those of normal spiral galaxies. In some cases, e.g. M82, starburst galaxies host both regions characterized by a SFE lower and higher than the mean of the sample. Furthermore, there are some galaxies for which the SFE is so high that the gas cycling timescale is $\\tau_{\\rm H_{2}} \\le 0.1$ Gyr, e.g. NGC3310."
    },
    "0201/astro-ph0201152_arXiv.txt": {
        "abstract": "Cold, dense clouds of gas have been proposed as baryonic candidates for the dark matter in Galactic haloes, and have also been invoked in the Galactic disc as an explanation for the excess faint sub-mm sources detected by SCUBA. Even if their dust-to-gas ratio is only a small percentage of that in conventional gas clouds, these dense systems would be opaque to visible radiation. This presents the possibility of detecting them by looking for occultations of background stars. We examine the possibility that the data sets of microlensing experiments searching for massive compact halo objects can also be used to search for occultation signatures by cold clouds. We compute the rate and timescale distribution of stellar transits by clouds in the Galactic disc and halo. We find that, for cloud parameters typically advocated by theoretical models, thousands of transit events should already exist within microlensing survey data sets. We examine the seasonal modulation in the rate caused by the Earth's orbital motion and find it provides an excellent probe of whether detected clouds are of disc or halo origin. ",
        "introduction": "\\label{intro}  Over the last decade several workers have suggested that the Galaxy may contain a large population of cold ($T \\la 10$~K), dense, low-mass clouds of molecular material. Several authors have argued that the mass contained in such clouds could be responsible for keeping galaxies' rotation curves flat beyond the edges of optical discs. \\citet{pfe94} argued that the clouds would be an integral part of a fractal interstellar medium and be confined to a thin disc. By contrast, \\citet{dep95}, \\citet{ger96}, \\citet{dra98}, and \\citet{wal98} considered the case of cold clouds that are distributed in an approximately spherical halo. De~Paolis et al, Draine and Gerhard \\& Silk argued for clouds of approximately solar mass, while Walker \\& Wardle pointed out that ionized winds from clouds with masses around a Jupiter mass ($\\mjup$) could explain the extreme scattering events that are frequently observed in the light curves of pulsars seen at low Galactic latitudes.  A major issue with these proposals that a significant mass is locked up in very small clouds, is to understand why the clouds do not collapse to form stars or degenerate objects such as brown dwarfs. \\citet{sci00} has argued that cosmic-ray heating may be able to balance molecular cooling in a stable way for dust-free clouds, whilst \\citet{law01} has presented similar arguments that cosmic-ray heating could balance cooling by dust emission in dusty clouds. Clouds that are an integral part of the ISM would be expected to be dusty, and radiate strongly at sub-mm wavelengths. \\citet{law01} has considered the possibility that a significant fraction of the reddest SCUBA sources could lie in the Galaxy. He concludes that for dusty clouds, masses between $\\sim 10\\mjup$ and $\\sim 10\\sm$ can be ruled out, and that a population that contributes significantly to the SCUBA counts will make a significant contribution to the mass of the Galaxy only if they have masses $\\sim M_{\\rm J}$. These objects would have to be confined to the Galactic plane.  The clouds would have temperatures $T\\sim 10$~K, and if they were in virial equilibrium, they would be $\\sim 20$~AU in diameter and totally opaque at optical wavelengths (and even have an optical depth of order unity at 1~mm).  Their number density on the sky has to be comparable to that, $\\sim 1000$~deg$^{-2}$, of unidentified SCUBA sources, so they would cover a fraction $\\sim 10^{-6}$ of the sky. Consequently, at any given time, of order one star in a million should be occulted by one of these clouds.  It should be feasible to detect such occultations in the databases that have been produced in searches for microlensing events \\citep[for a recent overview of microlensing results see][and references therein]{ker01}.  Here we investigate the rate at which stars would be occulted by opaque clouds.We argue in Section~\\ref{detectability} that halo and disc clouds are likely to have the characteristics required to produce observable transits. In Section~\\ref{theory} we derive general formulae for the rate and timescale distribution of transit events and in Section~\\ref{predictions} we apply them to a specific Galactic model. Currently available microlensing databases refer to a handful of lines of sight. The largest body of data refers to lines of sight in the general direction of the Galactic centre. Other extensive data sets refer to sight lines towards the Magellanic Clouds. Consequently, we concentrate on predictions for these particular directions. We also investigate means by which halo and disc cloud populations could be differentiated. ",
        "conclusions": "Cold opaque clouds can be detected by searching data sets already compiled by groups  searching  for gravitational microlensing events. Rather than the flux excess that is the signature of microlensing, cold cloud events are characterized by a transient dimming of background sources. Like microlensing, this signal should be non-periodic and one should expect events to involve a representative population of background source stars. In selecting microlensing events one problem is to reject variable stars which may mimic a microlensing signal. The equivalent ``background'' for cloud transit events is potentially much smaller since there are fewer classes of objects which involve a significant dimming of  flux. One such class, eclipsing binaries, is unlikely to present much of a problem. For short-period systems their true nature would be evident over the lifetime of the surveys. Binary systems with longer periods for which only one flux dropout is detected, should statistically comprise more massive stars, so we should expect the source stars of these ``events'' to be unrepresentative of the target stellar population as a whole, contrary to expectation for true cloud transit events. It is also likely that the light-curve of cloud transit events would have a form which generally would be inconsistent with that of eclipsing binaries.  The previous section demonstrates that the transit rate and timescales of cold clouds in the disc and halo is well within the range of detectability if they constitute a significant population and are of a Jupiter mass scale. Whilst the total rate depends only on the temperature of the clouds, virialized clouds more massive than $\\sim 0.05\\sm$ would have transit durations exceeding the present baseline of microlensing experiments.  The seasonal modulation of the rate provides a promising method to distinguish whether the clouds are in the disc or the halo. There are already strong arguments from the consideration of sub-mm sources \\citep{law01} that opaque cold clouds do not contribute much to the halo dark matter budget, though these arguments are sensitive to the precise dust properties of the clouds. Existing microlensing data sets represent a significant corpus of data which can provide an independent line of approach. Though originally conceived as a probe only of compact objects, microlensing surveys may also turn out to be one the most sensitive probes of diffuse cold clouds."
    },
    "0201/hep-ph0201184_arXiv.txt": {
        "abstract": "Using a form of modified dispersion relations derived in the context of quantum geometry, we investigate limits set by current observations on potential corrections to Lorentz invariance.  We use a phenomological model in which there are separate parameters for photons, leptons, and hadrons.  Constraints on these parameters are derived using thresholds for the processes of photon stability, photon absorption, vacuum \\v{C}erenkov radiation, pion stability, and the GZK cutoff.  Although the allowed region in parameter space is tightly constrained, non-vanishing corrections to Lorentz symmetry due to quantum geometry are consistent with current astrophysical observations. ",
        "introduction": "The quantum description of gravitation is arguably the largest gap in our understanding of fundamental physics.  In the last decade, a number of lines of research have offered new insights into the theory which will supersede quantum theory and general relativity.  For instance, several approaches predict that space is fundamentally discrete.  In one approach, eigenvalues of geometric observables have discrete spectra (see e.g. \\cite{lqgrev} for a review).  While it may not be surprising that the quantization of curved space yields quantized geometry, it is surprising that present day astronomical observations {\\em already} limit the extent of quantum geometry effects \\cite{kostelecky,AC,colglash,kifune,kluzniak,bert,aloisio,EMN99,EMN,ACP,steckglash,LJM,AP}.  These quantum geometry effects arise from an imprint of discrete underlying space on propagating modes.  Particles with ultra-high energies interact with structure on the smallest possible scales resulting in corrections to Lorentz symmetry.  Observations on the TeV scale offer an opportunity to test the extent of these quantum geometry induced Lorentz symmetry corrections.  Lorentz invariance may be the most tested symmetry in Nature.  Given the wealth of evidence in its support, it may seem obtuse to suggest that there may be corrections.  However, there are reasons to believe that this may not be an exact symmetry not the least of which is the fact that any test of Lorentz invariance (necessarily at finite energy) leaves an infinite parameter space untested.  Due to the non-compact nature of the Lorentz group, exact Lorentz symmetry is untestable.  There are other reasons to suspect exact Lorentz invariance at all energy scales.  For instance, ultraviolet divergences in quantum field theory point to new physics at high energies.  Despite these suggestions, only recently has work in quantum gravity yielded concrete proposals on how the symmetry might be modified \\cite{kostelecky,AC,GP,AM-TUneut,AM-TUphot,BACKG}.  In this paper we explore the consequences of modified dispersion relations which are motivated by a study of semiclassical states in loop quantum gravity by Alfaro, Morales-T\\'{e}cotl, and Urrutia \\cite{AM-TUneut,AM-TUphot}.  We study a model in which the usual dispersion relation of special relativity, $E^{2} = p^{2} + m^{2}$, is modified by leading order quantum geometry corrections of the form $\\kappa \\ell_{p} p^{3}$.  By the semiclassical analysis, the parameter $\\kappa$ is expected to be of order unity.  It determines the extent of the corrections to Lorentz symmetry while the Planck length, $\\ell_{p}$, sets the scale of the effects.  Due to the sensitivity of process thresholds, the modifications of particle dispersion relations are already tested by current astronomical observations.  What is more spectacular than merely limiting the possible extent of corrections to Lorentz symmetry is the proposal \\cite{acpion} that this framework may be robust enough to elegantly explain three incongruities, or paradoxes, between standard model predictions and observational results: (i) Cosmic rays are expected interact with the cosmic microwave background (CMB) producing pions and introducing an upper limit on the observed energy of particles of cosmological origin.  Known as the GZK cutoff \\cite{greisen,ZK} this upper limit has not been observed.  About twenty events at significantly higher energies have been reported \\cite{bird,takeda}.  Modified dispersion relations can raise the GZK cutoff \\cite{kifune}.  (ii) Ultra high energy photons of cosmic origin are also expected to interact with infra red background radiation.  According to some estimates for the background flux, photons of energy 10 TeV or more should not be seen due to background induced pair production \\cite{PM}.  However, higher energy events have been reported \\cite{tev1,tev2}.  Corrections to Lorentz invariance provides one explanation for this apparent paradox \\cite{kluzniak}.  (iii) Observations of longitudinal development in extensive air showers of high energy hadronic particles are apparently inconsistent with predictions \\cite{antonov}.  As proposed in Ref. \\cite{antonov}, one possible explanation is that high energy neutral pions become stable.  This may also be explained using modified dispersion relations \\cite{acpion}.  We calculate thresholds for processes involving photons, leptons and hadrons and, with observational limits, constrain the values for the $\\kappa$ parameters thereby confining the extent of quantum geometry corrections.  In more detail, in the next section we summarize the results of Alfaro {\\em et.  al.} \\cite{AM-TUphot,AM-TUneut}. Specifying only general properties of a semiclassical state, such as flatness above a characteristic scale $L$, the authors find that, in an analysis of particle propagation, photon \\cite{AM-TUphot} and neutrino \\cite{AM-TUneut} dispersion relations are modified.  In Section \\ref{process} we give a brief overview of threshold calculations before turning to a number of processes including: photon stability, photon non-absorption, vacuum \\v{C}erenkov radiation for electrons and protons, proton non-absorption, and pion stability.  We calculate constraints from the threshold calculations to investigate whether it is possible that observed effects may be accounted for by Lorentz symmetry corrections.  In \\ref{phot} we show that asymmetric momentum partitioning, first noticed by Liberati, Jacobson, and Mattingly \\cite{LJM}, dramatically affects the constraints on the dispersion relation modifications.  The results of Section \\ref{process} are summarized in Table \\ref{figure_table}.  Finally in Section \\ref{limits}, we apply the constraints together with current observations to limit the extent of potential Lorentz symmetry corrections.  We summarize the constraints in the final section and in Figures \\ref{egamma} and \\ref{pgamma}.  We find that present day observations tightly constrain -- but still leave open -- the possibility of Lorentz symmetry corrections of this form.  Particularly close to the present work is the paper by Jacobson, Liberati, and Mattingly \\cite{LJM} in which many of these results were summarized.  For the most part the present work agrees with this paper where the subject overlaps, although this work also includes new threshold calculations and constraints for proton vacuum \\v{C}erenkov radiation, the GZK threshold, and pion stability. ",
        "conclusions": "\\label{disc}  Beginning with a form of modified dispersion relations, we use exact energy-momentum conservation, to derive thresholds for particle processes.  The dispersion relations we consider are motivated by a class of semiclassical states of loop quantum gravity which are classical and flat above a characteristic length scale \\cite{AM-TUneut,AM-TUphot}.  We work in a model in which there are separate parameters for modifications of the dispersion relations for photons, leptons, and hadrons.  The process thresholds provide a sensitive test of the quantum geometry effects.  In some cases, such a photon decay and vacuum \\v{C}erenkov radiation, normally forbidden processes are activated at high energy.  In other cases, such a pion decay, processes are suppressed.  Current TeV-scale observations offer severe restrictions on the nature of these effects.  There exist allowed regions in both parameter spaces consistent with all of the above processes.  The results are summarized in Figures \\ref{egamma} and \\ref{pgamma} and in the last section.  In particular, including all of the processes the analysis shows that the hadronic parameter is effectively equal to the photon parameter and both of these are non-zero.  Further the lepton parameter is confined to a band in the third quadrant as shown in Fig. \\ref{egamma}.  \\begin{ack} We thank members of the Hamilton College Department of Physics, the Perimeter Institute, Hugo Morales-T\\'{e}cotl, and David Mattingly for helpful discussions during this work.  T.K. was supported, in part, by the Ralph E. Hansmann Science Students Support Fund of Hamilton College. \\end{ack}"
    },
    "0201/astro-ph0201236_arXiv.txt": {
        "abstract": "We are conducting a systematic lensing survey of X-ray luminous galaxy clusters at  $z\\sim0.2$ using the \\emph{Hubble  Space Telescope (HST)} and large ground-based telescopes.   We summarize initial results from our survey,  including a  measurement of the  inner slope of  the mass profile of  A\\,383, and a search for  gravitationally lensed Extremely Red Objects. ",
        "introduction": "Gravitational  lensing  by galaxy  clusters  is  a  powerful tool  for studying the distribution of mass in galaxy clusters at $z>0.2$ (e.g.\\ Kneib  et al.\\  1996 --  K96; Smith  et al.\\  2001a --  S01a)  and the properties of high-redshift ($z\\sim1$--6) galaxies (e.g. Smith et al.\\ 2002a --  S02a; Smail  et al.\\  2001; Ellis et  al.\\ 2001).   For this reason,  significant   effort  was  invested  during   the  1990's  in developing  the  lens   inversion  techniques  required  to  interpret robustly  cluster  lensing  observations  (K96).  During  this  period cluster lensing studies necessarily  concentrated on a small number of well studied clusters.  Our survey builds on the pioneering work of the 1990's and applies the K96  lens inversion  technique to  an objectively  selected  sample of clusters.  Ideally we would select  our cluster sample based on direct measurements of  their mass.  However,  in the absence of  large scale weak lensing surveys, we rely on X-ray luminosity as a crude indicator of cluster  mass for  the purpose of  sample selection.   We therefore select    12   X-ray    luminous   clusters    ($L_{\\rm   X}\\ge8\\times 10^{44}$\\,erg\\,s$^{-1}$, 0.1--2.4\\,keV) in  a narrow redshift slice at $z=0.17$--0.26, with  line of  sight reddening of  $E(B-V)\\le0.1$ from the XBACs sample (Ebeling et al.\\ 1996).  We  describe  the  first  three  published results  from  our  survey: detailed modeling  of the density  profile of A\\,383 (S01a);  a search for gravitationally  lensed Extremely  Red Objects (EROs)  (S02a); and near-infrared (NIR)  spectroscopy of a  dusty ERO uncovered  in S02a's survey (Smith et al.\\ 2001b -- S01b). ",
        "conclusions": "We are  conducting a  survey of 12  X-ray luminous galaxy  clusters at $z\\sim0.2$ with \\emph{HST} and  large ground-based telescopes.  One of these clusters,  A\\,383, contains  two radial arcs  which are  used to constrain the inner  slope of the cluster density  profile (S01a).  We find  that the cluster  scale dark  matter halo  appears to  possess a central cusp.  However, the logarithmic slope of this cusp is not well reproduced  by either  of the  rival theoretical  predictions  of dark matter density  profiles.  S02a also  exploit the magnifying  power of our cluster  lens sample to  construct a sample of  60 gravitationally lensed EROs  as faint  as $K\\sim22$.  Comparison  of our  observed ERO number  counts  with predictions  from  C00's  semi-analytic model  of galaxy formation reveals that  this models under-predicts the observed ERO number counts by an  order of magnitude.  NIR spectroscopy reveals that  one ERO in  our sample  is a  dusty starburst-Seyfert  galaxy at $z=1.05$ (S01b)."
    },
    "0201/astro-ph0201146_arXiv.txt": {
        "abstract": "s{ Cold Dark Matter (CDM) models of galaxy formation had been remarkably successful to explain a number of observations in the past decade. However, with both the theoretical modeling and the observations being improved, CDM models have been very recently shown to have excessive clustering on the sub-galactic scale.  Here I discuss a solution, based on our high-resolution numerical simulations, to this outstanding problem by considering Warm Dark Matter (WDM). Our results show that the over-clustering problem on sub-galactic scales can be overcome by WDM models, and all the advantages of CDM models are preserved by WDM models. Therefore, the WDM model will become an interesting alternative to the well-studied CDM models} ",
        "introduction": "Cold Dark Matter (CDM) models have been shown very successful to explain many observations of galaxies on scales of about one $\\mpc$ to a few hundred $\\mpc$. But such models probably predict over-clustering on smaller scales, as recent high-resolution simulations (Moore et al. 1999, Klypin et al. 1999, Jing \\& Suto 2000) have shown. The halo density profiles in these simulations are steeper than those inferred from the rotation curves of low surface brightness (LSB) galaxies, and there are too many sub-halos within galactic halos when compared to the observed number of satellite galaxies around the Milky Way. There is also additional evidence for such overclustering from, e.g.  the luminosity function of dwarf galaxies. Although some of these discrepancies may be resolved by introducing additional astrophysical processes (Bullock et al. 2000) and some others by properly interpreting the observations (van den Bosch \\& Swaters 2000), there are attempts to resolve the discrepancies by revisiting the assumption about the dark matter (DM). A list of the candidates for replacing CDM proposed since the summer of 1999 includes self-interacting DM, warm dark matter (WDM), repulsive DM, fuzzy DM, annihilating DM etc (see Dav\\`e et al. 2000 for references).  In this talk, I will present an extensive study for a warm dark matter model using high resolution N-body simulations. Our results will show that the WDM model is in good agreement with the observational data without resorting to not-well-understood astrophysical processes. A complete description of the study appeared in our recent paper submitted to the Astrophyical Journal (Jing 2000). ",
        "conclusions": ""
    },
    "0201/astro-ph0201370_arXiv.txt": {
        "abstract": "We have developed a new semi-analytic model for the formation and evolution of structure on galaxy, group and cluster scales. The model combines merger trees with a detailed, spatially resolved description of the dynamical evolution of halo substructure. It reproduces the results of numerical simulations of mergers and halo formation at a fraction of the computational expense. We use this model to study the formation of the dwarf galaxies in the Local Group and the structural components of the Milky Way. We find that reionization can explain the scarcity of Local Group dwarfs, although the reionization epoch is constrained to be high -- $z_{\\rm ri} \\ge 12$. We also find that disk disruption at recent times is rare, such that many galactic disks have an old, thin component. ",
        "introduction": "Substructure in CDM halos consists of hundreds of small, dense `subhalos'. The dynamical evolution of these satellites can be described using the analytic approach of Taylor \\& Babul (2001), which remains accurate to about 20\\% over many orbits. Combining this approach with extended Press-Schechter (EPS) merger trees, we have developed a new, semi-analytic (SA) model for the dynamical evolution of cluster or galaxy halos. This model predicts halo substructure identical to that seen in high-resolution numerical simulations. We have used the SA model to study the origin of the Local Group satellites, the stellar halo, and the thick and thin disk components of the Milky Way. ",
        "conclusions": ""
    },
    "0201/astro-ph0201416_arXiv.txt": {
        "abstract": "The \\hullac\\ atomic code is used to compute wavelengths and oscillator strengths for the 1s - $n$p inner-shell absorption lines in Li-like to F-like ions of neon, magnesium, aluminum, silicon, sulfur, argon, calcium, and iron. Many of these lines are expected to be observed in \\chandra\\ and \\xmm\\ high-resolution X-ray spectra of active galaxies. The new atomic data are incorporated in the ION code for spectral modeling of photoionized plasmas. The calculated spectra are subsequently compared with the spectrum of NGC~3783 and show very good agreement. The usefulness of these lines as diagnostics for the ionization state, column densities, and velocities in line-of-sight photoionized gas is called attention to. ",
        "introduction": "\\label{sec:intro} The launch of the \\chandra\\ and \\xmm\\ observatories along with the grating spectrometers on board has generated great interest in high resolution X-ray spectroscopy of a wide variety of astrophysical X-ray sources, in particular active galactic nuclei (\\agn). The high resolution grating observations of a handful of \\agn, e.g. NGC~5548  \\citep {kaastra00}, NGC~3783 \\citep {kaspi00, kaspi01}, \\mcg\\ \\citep {br01, lee01}, and \\iras\\ \\citep {sako01}, clearly show the presence of dozens of strong X-ray absorption lines originating from highly charged gas along the line of sight. While many of the absorption lines are due to transitions in H-like and He-like ions, some are clearly due to inner-shell transitions in lower ionization species. Perhaps the most prominent feature of this kind is the unresolved transition array (\\uta) of inner-shell $n$~= 2~to~3 (mainly 2p~- 3d) lines pertaining to various M-shell Fe ions identified in \\iras\\ \\citep {sako01}. The atomic data needed to model this \\uta\\ feature have recently been published in \\citet {behar01}. Another example is the inner-shell 1s - 2p feature of Li-like Si$^{11+}$  in NGC 3783 reported by \\citet {kaspi01}. The analogous 1s - 2p lines for Li-like O$^{5+}$ have been calculated by \\citet {pradhan00} using the R-matrix method and identified in \\mcg\\ by \\citet {lee01}. \\citet {sako02} have found evidence for additional 1s - $n$p lines arising from lower charge states of oxygen (O$^{4+}$ and O$^{3+}$) in \\mcg. Recently, \\citet {nahar01} published results for Li-like C$^{3+}$, O$^{5+}$, and Fe$^{23+}$ obtained using the same method as \\citet {pradhan00}.  Models like those presented in \\citet {kaspi01} predict that inner-shell transitions from the ground level of many more L-shell ions (Li-like to F-like) are indeed expected for X-ray illuminated photoionized plasmas. However, the available atomic data are few. The present paper aims at providing comprehensive results for the resonant 1s - $n$p transitions in Li-like through F-like ions of neon, magnesium, aluminum, silicon, sulfur, argon, calcium, and iron. For that purpose, we employ the \\hullac\\ code, which uses the parametric potential approximation. This method is by far more efficient than the previously used methods. For the present case of inner-shell transitions in highly charged ions, it is also expected to be very accurate. In the following, \\S2 describes the method of calculation and the new atomic data and \\S3 shows theoretical models for \\agn\\ spectra and a comparison with the spectrum of NGC~3783 obtained with the  \\hetgs\\ spectrometer on board \\chandra. ",
        "conclusions": "\\label{sec:disc} A complete set of wavelengths and oscillator strengths for inner-shell absorption lines are calculated for the entire L-shell of Ne, Mg, Al, Si, S, Ar, Ca, and Fe. This is a continuation of our ongoing study of inner-shell absorption features in \\agn\\ spectra. The usefulness of these atomic data for plasma diagnostics is demonstrated by comparison of spectral models with the \\hetgs\\ line-rich spectrum of NGC 3783. Good agreement is found between the model and the data, although for a full account of the X-ray absorber in NGC 3783, the \\hetgs\\ data clearly require a multi ionization-parameter model, which will be published separately."
    },
    "0201/astro-ph0201550_arXiv.txt": {
        "abstract": "We find convincing observational evidence to confirm the optical identification of the X-ray burster X1746-370 located in the globular cluster NGC\\thinspace6441. {\\it Chandra}/HRC-I imaging yields a much improved X-ray position for the source, which we show to be fully consistent with our rederived position of a UV-excess star, U1, in the same astrometric reference frame.  In addition, the smaller {\\it Chandra} X-ray error circle excludes the only other blue stars previously identified in the old {\\it Einstein} circle.  We have also obtained {\\it Hubble Space Telescope}/STIS time-resolved optical spectra of star U1.  Although there are no strong line features, the flux distribution demonstrates U1 to be unusually bright in the blue and faint in the red, consistent with earlier WFPC2 photometry.   More notably, the flux level of the continuum is seen to vary significantly compared to stars of similar brightness.  Indeed, the lightcurve can plausibly be fit by a 5.73 hr period sinusoid, which is the period of the recurring X-ray dips seen in this source. The presence of modulations in both wavelengths strengthens the case for an orbital origin, and therefore deepens the puzzle of the unusual energy independent X-ray dips. Lastly, we note that X1746-370 remains the longest period confirmed X-ray burster in a globular cluster, and the only one with a period typical of the galactic population as a whole. ",
        "introduction": "\\label{sect:intro} Globular clusters are expected to provide ideal environments for the formation of close binaries, with their high stellar densities and much enhanced rates of star interaction.   This is certainly the case for the X-ray bright interacting systems.  Almost from the advent of X-ray astronomy it has been known that the cluster population of luminous ($\\simgt10^{36}$\\ergsec) low-mass X-ray binaries (LMXBs) is $\\simgt$100 times enhanced (per unit stellar mass) relative to the galaxy as a whole  \\citep{clar75,katz75}.  Another curious fact is that 11 of these 13 LMXBs must contain neutron stars rather than black holes, as they exhibit type-I X-ray bursts (understood as thermonuclear runaway burning on the compact object's surface). Moreover, the sensitive X-ray imaging of the {\\it Chandra X-ray Observatory} has now revealed an equal number of probable quiescent systems in the few clusters examined to date \\citep{grin01a,grin01b,hein01,home01d,pool02,rutl02}, as well as the existence of {\\em two} persistent LMXBs in M15 \\citep{whit01}.  The study of globular cluster LMXBs benefits greatly from a multi-wavelength approach.  X-ray data, in general, only probe the vicinity of the central source, apart from the highest inclination systems where  material farther out can cross the line of sight.  The periods of four cluster LMXBs have been determined from eclipses by the donor star and/or from the periodic dips in their X-ray flux, understood as due to obscuration by vertically extended material near the edge of the accretion disc.  However, with the identification of optical/UV counterparts (in all but one case requiring the resolution of {\\it HST}), we can immediately begin to estimate the linear scale of a system from the $L_X/L_{opt}$ ratio, which has been shown to scale with disc area \\citep{vP94}.  Photometric monitoring has also proven effective in revealing variability on the binary period, whilst the optical spectra can in principle provide definitive corroboration of a counterpart and further useful diagnostics.  The X-ray burster  X1746-370, located in NGC\\thinspace6441, is one of the X-ray ``dippers''.  From a continuous {\\it EXOSAT} observation, \\citet{parm89} first observed dips and inferred an orbital periodicity, which was refined to 5.73$\\pm0.15$hr  by \\citet{sans93} using a more extensive {\\it Ginga} dataset.  Apart from a single deep (90\\% flux decrease) dip, which showed spectral hardening \\citep{jonk00}, all the observed dips have been shallow (\\til15\\%) and have shown no clear energy dependence. The apparently energy independent dipping is puzzling, implying that the obscuring material responsible for the electron scattering has metal abundance $\\simlt 0.01$ times solar, but this seems unlikely given the close to solar metallicity of the cluster as a whole \\citep{djor93}. The alternative explanations are: photoionization of the material, a number of varying spectral components conspiring together, or an extended X-ray source (i.e. an accretion disc corona).  However, the most recent broad-band spectroscopy of \\citet{parm99} using {\\it BeppoSAX} argues against any of these possibilities.  One might contend that the standard dipping interpretation itself could be erroneous-- and it is certainly true that none of the X-ray period determinations has been sufficiently precise to confirm the recurrent period as orbital in origin based on its stability. Additional progress has been made in the optical.  Using {\\it HST}/WFPC2 imaging data, \\citet{deut98} identified a variable, UV-excess star (designated U1) in the {\\it Einstein} X-ray error circle.  However, given the surprisingly large number of similar UV-bright stars in the cluster, there remained a possibility that U1 might be a chance superposition on the X-ray position. The {\\it a posteriori} probability of this coincidence was calculated to be \\til30\\% (based upon a 3\\arcsec\\ radius 90\\% confidence {\\it Einstein} error circle).  As part of a continuing program to probe the optical/UV counterparts to the luminous globular cluster X-ray sources, we have reexamined the optical position in the light of new {\\it Chandra} X-ray imaging data, and also obtained time-resolved {\\it HST}/STIS optical spectra of the candidate counterpart to the burster X1746-370. ",
        "conclusions": "The new {\\it Chandra} X-ray and optical positions for X1746-370 and its proposed counterpart U1 are summarized in table~\\ref{tab:posns}, and illustrated in figure~\\ref{fig:posns}.  The positional agreement is excellent and well-within the (small) estimated radial uncertainties, providing strong support for the optical identification. Following the {\\it a posteriori} probability estimate of \\citep{deut98}, but with a new 90\\% confidence radius of \\decsec{1}{1} for the {\\it Chandra} error circle, the area enclosed is 7 times smaller than for {\\it Einstein}, and hence only a $\\sim$4\\% probability remains that we have chance alignment of an unrelated UV-excess cluster star with the X-ray position.  Moreover, our new {\\it HST}/STIS spectra confirm that star U1 is unusually blue. But more conclusively, the temporal coverage afforded by our time-resolved spectra are sufficient to confirm the optical variability of the source, which appears periodic and can be well fitted with a sinusoid constrained to the 5.73 hr X-ray period.  The firm identification of the optical counterpart does have further implications for our view of the nature of the source, including the issue of the unusual energy independent X-ray dips. First, the most likely origin for the sinusoidal optical modulation on the X-ray period is the varying contribution from the bright X-ray heated face of the donor.  Hence, its possible detection would not only confirm the correct identification, but also support the standard picture that the recurring dips are related to periodic obscuration on the orbital period.  Second, as previously noted by  \\citet{parm99}, the faintness of the optical star implies that $L_X/L_{opt}\\sim1000$, typical of LMXBs in which we directly observe the central source, and therefore consistent with the detection of bursts from this source. In particular, this suggests that it is unlikely that the dips are primarily due to obscuration of an extended accretion disc corona (ADC); in the classical ADC sources where only scattered X-rays are visible, $L_X/L_{opt}\\sim20$.  Our optical spectral data also argue against such an ADC interpretation. Considering all LMXBs with comparable periods (and hence disc sizes), most show characteristic emission lines in this spectral region at H$\\beta$, He{\\scs II} $\\lambda$4686 and the Bowen C{\\scs III}/N{\\scs III} blend $\\lambda$4640; moreover, all three known classical ADC sources exhibit strong emission at one or more of these lines \\citep[see e.g.][]{vP95}.  Hence, the lack of strong line emission in X1746-370 is unlike any of these ADC systems, and is even unusual compared to the field LMXBs in general.  However, we note that our {\\it HST} spectra of the globular cluster LMXBs in NGC\\thinspace6624, and 6712 \\citep{deut98T} are similar to X1746-370 in NGC\\thinspace6441 insofar as they are also very blue/UV, but largely featureless at modest signal-to-noise and resolution.  Within the context of the diverse nature of the cluster LMXB population, X1746-370 might at first glance be considered a rather average system.  In terms of period it lies midway between the three with ultra-short periods (P$<$ 1 hr) in NGC\\thinspace6624, 1851 and 6712 \\citep[see e.g.][and references therein]{home01c}, and the two long period systems in Terzan 6 and M15 (AC211) \\citep[P=12.4 hr and 17.1 hr respectively;][]{intZ00,ilov93}. However, it is the longest period confirmed burster, since the Terzan 6 LMXB has not been seen to burst, and we now know that AC211 (M15-X1) is almost certainly an ADC source and M15-X2, a source without a known orbital period, is the probable burster there \\citep{whit01}. The recent results on NGC\\thinspace6652 \\citep{hein01} indicate that its burster is perhaps the most similar, with a longest possible period of 4.4 hr, though 0.92 hr also fits the available data.  If this shorter period in NGC\\thinspace6652 does turn out to be correct, of the five bursters with determined or well-constrained periods, four would be double-degenerate ultra-compact systems and only X1746-370 would be similar to a typical galactic burster. \\citet{deut00} have already commented on this prevalence of very exotic systems in globular clusters.  It would appear that the unique formation/evolution processes at work in globular cluster cores \\citep[see e.g.][for a review]{hut92}-- e.g. the tidal capture and exchange encounter mechanisms and subsequent stellar interactions, leading to the hardening of already hard binaries -- may have led to an enhancement of the ultra-compact LMXBs at the expense of the wider systems like X1746-370. In any case, the (growing) population of LMXBs may serve as important tracers of the stellar dynamics and evolution within globular clusters."
    },
    "0201/astro-ph0201285_arXiv.txt": {
        "abstract": "Despite great observational and theoretical effort, the burst progenitor is still a mysterious object. It is generally accepted that one of the best ways to unveil its nature is the study of the properties of the close environment in which the explosion takes place. We discuss the potentiality and feasibility of time resolved X--ray spectroscopy, focusing on the prompt $\\gamma$-ray phase. We show that the study of absorption features (or continuum absorption) can reveal the radial structure of the close environment, unaccessible with different techniques. We discuss the detection of absorption  in the prompt and afterglow spectra of several bursts, showing how these are consistent with gamma-ray bursts taking place in dense regions. In particular, we show that the radius and density of the surrounding cloud can be measured through the evolution of the column density in the prompt burst phase. The derived cloud properties are similar to those of the star forming cocoons and globules within molecular clouds. We conclude that the burst are likely associated with the final evolutionary stages of massive stars. ",
        "introduction": "It is widely believed that a good way to understand which is the progenitor of GRBs is by analyzing the properties of the interstellar medium that surrounds the explosion. This is because the fireball early self--similar evolution erases all the traces of its initial condition and hence any pre and during--explosion signature.  The three main classes of burst progenitor can be, in principle, easily distinguished by mean of the properties of their environment. If the burst are due to the merger event of binary coalescing systems of neutron stars \\citep{eic89} they are expected to take place in a uniform low density ($n\\sim0.1-10$~cm$^{-3}$) intergalactic medium. This is due to the fact that the binary system has a long life ($\\sim 10^9$~y) and a high proper motion ($v\\sim100-1000$~km/s) and can travel out of the original birth place before the merging event (but see Perna et al., this volume)  If the explosion of a GRB is coincident with the explosion of a massive rotating star (hypernov\\ae~or collapsar, Woosley, this volume), it has to be surrounded primarily by the pre--explosion stellar wind. This wind will then impact on the molecular cloud in which the star was born, with a shock contact (terminal) discontinuity \\citep{ram01}.  Finally, the bursts may be associated to supernova explosions but with some delay (see the supranova model \\citep{vie98}). In this case the burst should explode in an evacuated cavity, surrounded by a supernova shell and eventually by a molecular cloud medium.  These three radial profiles of the ambient media surrounding GRBs are sketched in the left panel of Fig. 1, where the solid line represents compact mergers and the dashed and dotted lines represent hypernovae and supranovae, respectively.  In principle the density profile can be traced by modelling the afterglow light curves and spectra. One must however be aware that most afterglow data are taken between half a day and several months after the burst explosion. As it is shown in the right panel of Fig. 1, in this period of time the fireball, no matter the progenitor model, runs through a uniform medium, the only difference being the normalization (probably the most uncertain of all the model parameters). The morphological difference of the left panel is then impossible to reconstruct with present day measurements.  There are several alternatives in order to measure the close environment density and structure. One is to be fast. In principle, if one can have (as we will have in the Swift era) detailed early time light curves, the whole radial density structure can be measured. However, it is likely that the emission mechanism of the afterglow gets more complicated as we approach the explosion site: reverse shock emission, late injection of energy from the inner engine and the superposition with the radiation from internal shocks will probably make the modeling of early time afterglows a delicate issue.  \\begin{figure} \\includegraphics[width=\\textwidth]{envi} \\caption{Radial density profiles for different GRB progenitors. the solid line represent compact mergers and the dashed and dotted lines represent hypernovae and supranovae, respectively. The left panel shows the pre--explosion setup, while in the right panel only the range of radii that the fireball travels between an observer time $t=12$~hours and $t=2$~months is highlighted.} \\end{figure}  \\begin{figure} \\includegraphics[width=0.48\\textwidth]{simule} \\caption{Opacity in the range [2-10]~keV for a cloud with solar metallicity, $R=3\\times10^{18}$~cm and initial column density $N_H(0)=3\\times10^{21}$~cm$^{-2}$. In the main panel, from top to bottom, we plot the absorption at times $t=0$, 10, 40, 80, 120 and 160 seconds. In the inset, the column density is shown as a function of time. Filled dots mark the column densities corresponding to the spectra plotted in the main panel.} \\end{figure}  An alternative is to look for echoes, i.e. photons that, initially emitted at large angles with respect to the line of sight, are scattered in the direction of the observer. Dust \\citep{esi00}, Compton \\citep{mad00} and iron line \\citep{laz99} echoes have been proposed. Only iron lines have been securely observed, to date \\citep{pir00}. The modelling of echoes presents two difficulties: first, if GRB fireballs are highly collimated, there is little room for echoes. Secondly, it is difficult to disentangle photons scattered by a large angle at small distance from the burst site from photons scattered at small angles at a larger distance from the progenitor.  We here propose and analyze a method, based on prompt time--resolved X--ray spectroscopy of the burst photons, which is unbiased and rely on very well known physics. The propagation of the photons in the ambient medium will in fact imprint absorption features on the soft X--ray spectra. These features will become less deep as the ionization front expands, allowing us to measure the density and the radial profile of the surrounding material. ",
        "conclusions": "We have shown that time resolved X--ray spectroscopy of the early phases of GRB emission can give us informations on the density and radial structure of the surrounding material. Given the capabilities of present days instrumentations, this can be effectively done in case of fairly dense and compact regions. If we impose that the column density must not be negligible after 1 second of observation and that it must decrease by a factor of two after 100 seconds of GRB emission, we find that, for a uniform absorbing cloud, positive detections of $N_H$ variations should be performed if a GRB is surrounded by a cloud with size and column density marked with the gray shading in Fig. 6.  We compared these cloud properties with typical properties of molecular clouds and their overdense regions in our Galaxy. We find that if GRBs take place in random locations inside molecular clouds, their X--ray early spectra should show no sign of photoionization absorption, since all the material is ionized on a time scale of less than one second. On the other hand, massive stars are thought to be born inside overdense and compact regions within molecular clouds. If GRBs are associated with these regions, evolution of the X--ray absorbing column should be detectable. In particular, the variable absorption observed in the spectra of GRB~980329 and GRB~780506 can be explained if they were located in regions with properties close to those of Bok globules.  \\begin{figure} \\includegraphics[width=0.48\\textwidth]{molclu} \\caption{Comparison between the radius and column densities inferred and the properties of Galactic molecular clouds and some hierarchical structures embedded in them. The gray and line shaded areas are where variable column is observable with GRB measurements. Circles show radii and average column densities of a sample of Galactic molecular clouds. Asterisks refer to Bok globules, diamonds refer to dense cores in the Taurus molecular cloud and triangles refer to massive cloud cores. The two filled dots are the best fit values to the $N_H$ measurements of GRB~780506 and GRB~980329.} \\end{figure}   \\doingARLO["
    },
    "0201/astro-ph0201440_arXiv.txt": {
        "abstract": "Understanding a background is crucial in particular for a study of low surface brightness objects. In order to establish the background subtraction method, we have studied properties of the EPIC background.  Count rates of the background vary violently by two order of magnitude at the maximum, while during the most quiet period, these are stable within 8 \\% at a 1 $\\sigma$ level. The overall spectrum is dominated by particle events above 5 keV, and its spatial variation is also found.  The long-term variation of the background is also investigated with CAL CLOSED data, which is the data of calibration source with filter closed. The average background count rate decreased by 20 \\% from March 2000 to January 2001, but it regained in February 2001.  For the modeling of the background spectrum, we investigate correlations between the 2-10 keV count rate and some characteristic parameters. The PN background shows a good correlation with some parameters. On the other hands, the MOS background does not shows a clear correlation. Further investigation is needed for the MOS background.  Our final goal is to establish a method to predict the background, for which these results will be reflected in the background generator. ",
        "introduction": "In order to study properties of the background, we used data from Science PV phase data operated in the full frame mode with thin filter. All the data are shown in table \\ref{tbl:hkatayama-WA3_tbl1} and table \\ref{tbl:hkatayama-WA3_tbl2}. The number of data sets is 8 for PN and 7 for MOS.  Figure \\ref{fig:hkatayama-WA3_fig1} shows a light curve of PN from one of the Lockman hole observations. In order to exclude the flare events, time periods where the count rate deviates from the mean value during quiescent periods by $\\pm2 \\sigma$ are excluded. Celestial sources and noisy columns are also excluded from the data. The average exposure time thus remained after excluding the flare events is 18 ks for PN and 26 ks for MOS.    \\begin{table} \\begin{center} \\begin{tabular}{l l l}\\hline Obs ID\t                & Exposure &   Object   \\\\ \\hline\\hline 0063\\_0123100201\\_PNS001\t& 11411\t  & MS0737.9+7441 \\\\ 0070\\_0123700101\\_PNS003\t& 27695\t  & Lockman Hole \\\\ 0071\\_0123700201\\_PNS003\t& 21009\t  & Lockman Hole \\\\ 0073\\_0123700401\\_PNS003\t& 9977\t  & Lockman Hole \\\\ 0078\\_0124100101\\_PNS003\t& 21403\t  & RXJ0720.4-3125 \\\\ 0081\\_0123701001\\_PNS003\t& 17569\t  & Lockman Hole \\\\ 0082\\_0124900101\\_PNS003\t& 17665\t  & MS1229.2+6430 \\\\ 0181\\_0098810101\\_PNS003\t& 18529\t  & WW Hor \\\\ \\hline \\end{tabular} \\caption{PN data summary} \\label{tbl:hkatayama-WA3_tbl1} \\end{center} \\begin{center} \\begin{tabular}{l l l}\\hline Obs ID\t &              Exposure\t& Object \\\\ \\hline\\hline 0070\\_0123700101\\_M1S001\t& 30552\t  & Lockman Hole \\\\ 0071\\_0123700201\\_M1S001\t& 33269\t  & Lockman Hole \\\\ 0078\\_0124100101\\_M1S001\t& 28794\t  & RXJ0720.4-3125 \\\\ 0081\\_0123701001\\_M1S001\t& 21358\t  & Lockman Hole \\\\ 0100\\_0123701001\\_M1S001\t& 21258\t  & Lockman Hole \\\\ 0071\\_0123700201\\_M2S002\t& 26465\t  & Lockman Hole \\\\ 0181\\_0098810101\\_M2S002\t& 21785\t  & WW Hor \\\\ \\hline \\end{tabular} \\caption{MOS data summary} \\label{tbl:hkatayama-WA3_tbl2} \\end{center} \\end{table}   \\begin{figure}[htbp] \\begin{center} \\psbox[xsize=0.2#1,ysize=0.2#1]{hkatayama-WA3_fig1.eps} \\caption{Light curve of Lockman hole observation. } \\label{fig:hkatayama-WA3_fig1} \\end{center} \\end{figure} ",
        "conclusions": ""
    },
    "0201/astro-ph0201489_arXiv.txt": {
        "abstract": "We present thermal infrared images of the bipolar nebula surrounding $\\eta$ Carinae at six wavelengths from 4.8 to 24.5 $\\micron$. These were obtained with the MIRAC3 camera system at the Magellan Observatory.  Our images reveal new intricate structure in the bright core of the nebula, allowing us to re-evaluate interpretations of morphology seen in images with lower resolution.  Complex structures in the core might not arise from a pair of overlapping rings or a cool (110 K) and very massive dust torus, as has been suggested recently. Instead, it seems more likely that the arcs and compact knots comprise a warm ($\\sim$350 K) disrupted torus at the intersection of the larger polar lobes.  Some of the arcs appear to break out of the inner core region, and may be associated with equatorial features seen in optical images.  The torus could have been disrupted by a post-eruption stellar wind, or by ejecta from the Great Eruption itself if the torus existed before that event.  Kinematic data are required to rule out either possibility. ",
        "introduction": "An elegant expanding bipolar nebula known as the Homunculus surrounds the persistently peculiar star $\\eta$ Carinae.  The nebula contains several $M_{\\odot}$ of material, most of which was ejected during a major eruption about 160 years ago (see Davidson \\& Humphreys 1997). {\\it Hubble Space Telescope (HST)} images at optical wavelengths (Morse et al.\\ 1998) show interesting detailed structure in the polar lobes, as well as a ragged equatorial debris disk.  Dust in the Homunculus absorbs most of the star's UV and optical flux, and then re-emits roughly $4 \\ \\times \\ 10^6 \\ L_{\\odot}$ at infrared (IR) wavelengths, making $\\eta$ Car one of the brightest sources in the sky at 10 $\\micron$ (Westphal \\& Neugebauer 1969), despite its distance of 2.3 kpc (Davidson \\& Humphreys 1997).  IR radiation is an especially useful way to study $\\eta$ Car because it reveals structures {\\it inside} the Homunculus that are mostly obscured at shorter wavelengths.  Numerous investigations at IR wavelengths have shown a bright, elongated, and multiple-peaked core a few arcseconds across (e.g., Hyland et al.\\ 1979; Mitchell et al.\\ 1983, Hackwell et al.\\ 1986; C.H.\\ Smith et al.\\ 1995; Rigaut \\& Gehring 1995; Smith et al.\\ 1998; Polomski et al.\\ 1999; Smith \\& Gehrz 2000).  These features are usually interpreted as an inclined and limb-brightened circumstellar torus or disk.  The sharpest mid-IR image of $\\eta$ Car so far was presented by C.H.\\ Smith et al.\\ (1995), who used a special processing technique to produce a remarkable 12.5 $\\micron$ image showing the detailed spatial structure of the core, consisting of several knots and loop structures.  Smith et al.\\ (1998) and Polomski et al.\\ (1999) discussed multi-wavelength IR array images of $\\eta$ Car, and summarized this object's unique thermal-IR emission properties.  A controversey has recently arisen regarding the nature of the bright structures in the core of the Homunculus.  Several previous observers had described the torus and showed that it had a color temperature higher than 250 K at wavelengths near 10 $\\micron$ (e.g., Hackwell et al.\\ 1986; C.H.\\ Smith et al.\\ 1995; Smith et al.\\ 1998; Polomski et al.\\ 1999).  However, Morris et al.\\ (1999) recently claimed to have discovered this feature, and their interpretation of a large-aperture 2 to 200 $\\micron$ spectrum obtained with the {\\it Infrared Space Observatory (ISO)} led them to associate the torus with much cooler dust at 110 K.  They proposed that this cool torus contained $15 \\ M_{\\odot}$ of material, and was formed when a companion star in a close binary system stripped the normal composition envelope off the primary star before the Great Eruption.  They further suggested that the ejection of this massive torus caused the Great Eruption (by increasing the star's $L / M$ ratio) and was directly responsible for the bipolar shape of the Homunculus Nebula.  However, Davidson \\& Smith (2000) showed that the hypothesized compact torus inside the Homunculus would be too small to radiate the required luminosity with a brightness temperature of only 110 K. Hony et al.\\ (2001) proposed instead that the inner dust was warmer, as already shown by earlier studies, but they interpreted structure seen in their images as a pair of overlapping rings with a geometry analogous to those around SN 1987a (e.g., Burrows et al.\\ 1995).  These hypothetical rings have a different polar axis than the Homunculus --- so Hony et al.\\ (2001) made the extraordinary claim that an interaction between three bodies had caused the orbital axis to change orientation by more than one radian during or after the Great Eruption in the mid-19th century when the Homunculus was ejected.  This interpretation seems unlikely for several reasons, and here we present new IR data showing that interpreting the IR features inside the Homunculus as a pair of rings is less straightforward than Hony et al.\\ suggest.  Instead these features could comprise an equatorial disk of material as previously proposed, which appears significantly disrupted when observed at high spatial resolution.  In \\S 2 we present our new thermal-IR images, and in \\S 3 we discuss the IR structures observed in these high-resolution images, as well as their consequences for previous interpretations of $\\eta$ Car's IR radiation.  This will be a brief and descriptive analysis of the IR structures in the core of the nebula; a more thorough quantitative analysis of these data will follow in a later paper. ",
        "conclusions": ""
    },
    "0201/astro-ph0201395_arXiv.txt": {
        "abstract": "I consider the thermal bremsstrahlung emission from hot accretion flows (Bondi/ADAFs), taking into account the finite size of the observing telescope's beam ($R_{beam}$) relative to the Bondi accretion radius ($R_{A}$). For $R_{beam} \\gg R_{A}$ soft X-ray emission from the hot interstellar medium surrounding the black hole dominates the observed emission while for $R_{beam} \\ll R_{A}$ hard X-ray emission from the accretion flow dominates. I apply these models to {\\it Chandra} observations of the Galactic Center, for which $R_{beam} \\approx R_{A}$.  I argue that bremsstrahlung emission accounts for most of the ``quiescent'' (non-flaring) flux observed by {\\it Chandra} from Sgr A*; this emission is spatially extended on scales $\\sim R_{A} \\sim 1''$ and has a relatively soft spectrum, as is observed.  If accretion onto the central black hole proceeds via a Bondi or ADAF flow, a hard X-ray power law should be present in deeper observations with a flux $\\sim 1/3 $ of the soft X-ray flux; nondetection of this hard X-ray component would argue against ADAF/Bondi models. I briefly discuss the application of these results to other low-luminosity AGN.   \\  \\noindent {\\it Subject Headings:} Galaxy: center --- accretion, accretion disks ",
        "introduction": "Spherical Bondi accretion predicts that the interstellar medium around a black hole should be gravitationally captured on scales of $R_A \\approx GM/c^2_s$, where $M$ is the mass of the black hole and $c_s$ is the sound speed of gas in the vicinity of $R_A$ (e.g., Bondi \\& Hoyle 1944; Shvartsman 1971).  The same is true for hot accretion flow models that include dynamically important angular momentum, such as advection-dominated accretion flows (ADAFs; e.g., Ichimaru 1977; Rees et al. 1982; Narayan \\& Yi 1994) and its variants (e.g., Blandford \\& Begelman 1999).  For radii $\\lsim R_A$ the captured gas accretes onto the central black hole and the dynamics is determined by accretion physics rather than interstellar medium physics.  At a minimum, the gas in Bondi and ADAF models emits thermal bremsstrahlung emission from radii $\\lsim R_A$ (e.g., Narayan et al. 1998; Di Matteo et al. 1999, 2000).  Additional X-ray emission from synchrotron or inverse Compton processes in the accretion flow or jet can also be present depending on the accretion rate and the efficiency of electron acceleration (e.g., Narayan et al. 1998; Markoff et al. 2001; Yuan et al. 2002; Narayan 2002).  For sufficiently low-luminosity systems bremsstrahlung emission may dominate over other emission processes in the X-ray band; e.g., Narayan et al. (1999; see Fig. 6) found that bremsstrahlung dominated for $L_X \\lsim 10^{-8} L_{\\rm Edd}$.  This range of luminosities is now routinely probed by {\\it Chandra} observations (e.g., Ho et al. 2001).  In particular, for the Galactic Center $L_X \\sim 10^{-11} L_{\\rm Edd}$ (Baganoff et al. 2002) and so it is {\\it a priori} plausible that bremsstrahlung contributes significantly to the observed emission.  For a given system, the relative contribution of the accretion flow ($R \\lsim R_A$) and the ambient medium ($R \\gsim R_A$) to the thermal bremsstrahlung emission depends on the size of the observing telescope's beam in units of the Bondi accretion radius.  For $R_{beam} \\gg R_A$ the ambient medium dominates the observed emission while for $R_{beam} \\ll R_A$ the accretion flow does.  The interpretation of observed data therefore depends sensitively on $R_{beam}/R_A$.  For ambient temperatures of $\\approx 1$ keV, $R_A \\approx 0.07$ pc for the $2.6 \\times 10^6 M_\\odot$ black hole at the Galactic Center and $\\sim 30$ pc for the $\\sim 10^9 M_\\odot$ black holes in massive elliptical galaxies such as M87, NGC 4472, and NGC 1399 (for the Galactic Center, see Genzel et al.  1997 and Ghez et al. 1998 for a black hole mass estimate and Baganoff et al. 2002 for a temperature measurement; for elliptical galaxies, see, e.g., Gebhardt et al. 2000 and Ferrarese \\& Merritt 2000 for black hole mass estimates and Loewenstein et al. 2001 for central temperature measurements).  For comparison, the $\\approx 1\"$ angular resolution of the {\\it Chandra} X-ray Observatory corresponds to a distance of $\\approx 0.04$ pc at the Galactic Center and $\\approx 85$ pc in nearby X-ray clusters such as Virgo and Fornax.  {\\it Chandra} observations thus probe length scales comparable to the Bondi accretion radius for a number of supermassive black holes.  In this paper I present models for the bremsstrahlung emission from hot accretion flows that can be applied to {\\it Chandra} observations with $R_{beam} \\sim R_A$.  This paper is organized as follows.  In the next section (\\S2) I show X-ray spectra for bremsstrahlung emission from Bondi accretion flows that explicitly account for the finite size of the observing telescope's beam.  I then compare these predictions with {\\it Chandra} observations of Sgr A* at the Galactic Center (\\S3). In \\S4 I conclude and discuss additional applications of these results. ",
        "conclusions": "In this paper, I have presented model X-ray spectra for bremsstrahlung emission from Bondi accretion flows that explicitly include the contribution from hot ambient gas around the black hole (Fig. 2). These spectra are useful for interpreting {\\it Chandra} observations of very low-luminosity AGN since theoretical models suggest that bremsstrahlung may dominate over other emission processes for $L \\ll L_{Edd}$ (perhaps $L_X \\lsim 10^{-8} L_{\\rm Edd}$). Care must be taken in interpreting observations of these systems because the ambient gas around the black hole contributes significantly to the observed emission if $R_{beam} \\gsim R_A$ (see Fig. 2), as is the case for most known systems even given {\\it Chandra's} excellent angular resolution.  I have applied these results to {\\it Chandra} observations of Sgr A* at the Galactic Center. I propose that, excluding the large X-ray flare, bremsstrahlung emission from gas in the vicinity of the Bondi accretion radius dominates the quiescent flux observed from Sgr A*. This model explains why the observed spectrum is quite soft, why the quiescent flux is relatively constant (i.e., independent of time), and why the source is spatially extended (\\S3). In this interpretation the quiescent emission from Sgr A* does not presently constrain accretion flow models because the emission arises from gas at $\\sim R_A$, i.e., from the ``transition region'' between the ambient medium and the accretion flow.  If, however, accretion onto the black hole proceeds via a Bondi or ADAF flow, a hard X-ray power law should be detectable in deeper {\\it Chandra} observations. The flux in this power law comes from gas at $\\ll R_A$ and should be $\\sim 1/3$ of the soft thermal flux from gas at $\\sim R_A$ (Fig. 2a; middle curve).  It is important to stress that {\\it Chandra} observations of the Galactic Center directly determine the density and temperature of gas at $1.5'' \\approx R_A$ ($\\approx 100$ cm$^{-3}$ and $\\approx 2$ keV, respectively) and at $10'' \\approx 7 R_A$ ($\\approx 20$ cm$^{-3}$ and $\\approx 1$ keV, respectively).  These boundary conditions strongly constrain Bondi accretion models and it appears difficult to avoid the hard X-ray power law seen in Figure 2 (if Bondi models are correct).  One caveat is that the models presented here assume that all of the inflowing gas is subsonic at large radii ($\\gsim 1''$).  Since Sgr A* is believed to be fed by stellar winds with velocities $\\sim 300-1000$ km s$^{-1}$ (Najarro et al. 1997), I have effectively assumed that most of the stellar winds have shocked outside $\\approx 1''$.  For the X-ray emitting gas observed by {\\it Chandra} this is probably reasonable since the observed temperature of several keV is comparable to that expected from the shocked stellar winds.  There could, however, be a component of cold inflowing gas that is not accounted for in the models presented here.  Advection-dominated accretion flow models require an additional boundary condition on top of the density and temperature needed in Bondi models, namely the rotation rate of the gas at $\\approx R_A$. It would be interesting to explore the predicted spectra from ADAF models as a function of this rotation rate (which is rather uncertain).  On physical grounds ADAF models should be very similar to the Bondi models shown here as long as the viscous time at $\\sim R_A$ is shorter than the cooling time; this is easily satisfied for the Galactic Center provided the dimensionless viscosity $\\alpha$ is $\\gsim 10^{-3}$.  My preliminary calculations of ADAF models accreting from an ambient medium support this conclusion.  Although the presence of an underlying hard X-ray power law at roughly the level predicted here appears relatively robust within the context of ADAF and Bondi models, it does depend sensitively on the structure of the accreting gas at $R \\lsim R_A$. Moreover, recent theoretical work suggests that that the dynamics of radiatively inefficient accretion flows may be quite different from that predicted by Bondi and ADAF models (e.g., Blandford \\& Begelman 1999; Stone et al. 1999; Igumenshchev \\& Abramowicz 2000; Narayan et al. 2000; Quataert \\& Gruzinov 2000; Igumenshchev \\& Narayan 2002; Hawley \\& Balbus 2002). For a parameterized density profile of the form $\\rho \\propto R^{-3/2 + p}$ ($0 < p < 1$), large values of $p \\sim 1/2-1$ are favored over the Bondi/ADAF value of $p = 0$.  It is unclear whether these modifications to the structure of the accretion flow occur as far out as $\\sim R_A$ or whether they are confined to regions closer to the black hole.  For example, most of the physics that leads to significant deviations from ADAF/Bondi models requires dynamically significant angular momentum (see, however, Igumenshchev \\& Narayan 2002).  It is unclear whether this is appropriate near $\\sim R_A$, where most of the detectable bremsstrahlung emission originates.  Thus, even if the flow structure is very different in the vicinity of the black hole, Bondi models may be applicable near $\\sim R_A$.\\footnote{This is a subtle issue because, even if the angular momentum barrier formally lies at a radius $\\ll R_A$, angular momentum can still influence the flow structure out to $\\sim R_A$.  This is because the flow is causally connected from near the horizon out to large radii.}  If the flow structure is significantly modified in the vicinity of $R_A$, one does not expect to see a hard X-ray power law analogous to the Bondi/ADAF results in Figure 2. The reason is that the bremsstrahlung spectrum from a $\\rho \\propto R^{-3/2 + p}$ density profile is $\\nu L_\\nu \\propto \\nu^{1/2 - 2p}$ (Quataert \\& Narayan 1999). For $p = 0$ (Bondi/ADAF), a hard X-ray power law is present (Fig. 2) while for $p \\sim 1/2-1$, the spectrum is soft and so there should be very little hard X-ray emission. Deep {\\it Chandra} observations of the quiescent emission from the Galactic Center can therefore shed important light on the structure of the accretion flow onto the central black hole.  Interpreting the data may be complicated if, as is plausible, SSC and/or synchrotron emission contribute to the quiescent flux at some level.  The bremsstrahlung contribution can be isolated by focusing on the least variable segments of the observed data.  The strength of thermal X-ray lines also provides a constraint on synchrotron or SSC contributions to the quiescent emission (Narayan \\& Raymond 1999).  The constraints on an underlying hard X-ray power law can be significantly tightened if there are sufficient counts to fully utilize {\\it Chandra's} resolution and extract spectra from a $\\approx 0.5''$ region around Sgr A*.  As Figure 2a shows, most of the soft thermal emission would then be resolved out (since $R_{beam} \\approx 0.3 R_A$).  It might even be possible to construct a radial surface brightness profile of the outer parts of the accretion flow to compare with theoretical models (e.g., Quataert \\& Narayan 2000; \\\"Ozel \\& Di Matteo 2001).  For example, in ADAF/Bondi models, the spectrum should harden significantly as $R_{beam}/R_A$ decreases from 1 to 0.3 (Fig. 2).  \\subsection{Additional Applications}  The results in this paper are also useful for interpreting {\\it Chandra} observations of massive black holes in early type galaxies. In such systems the central black hole may accrete the hot ambient interstellar medium of the host galaxy (e.g., Fabian \\& Rees 1995; Di Matteo et al. 1999).  Loewenstein et al. (2001) reported the nondetection of a central hard X-ray point source in NGC 1399, NGC 4472, and NGC 4636.  Given black hole mass estimates from the $M_{bh}-\\sigma$ relation (Gebhardt et al. 2000; Ferrarese \\& Merritt 2000), $R_{beam} \\approx 3, 4$, and $20 \\ R_A$ for these systems.  The ambient ISM in elliptical galaxies is typically stratified as $\\rho \\propto R^{-1}$ so the results in Figure 2a can be used to estimate the expected hard X-ray bremsstrahlung emission from a Bondi or ADAF accretion flow.  I find that the hard X-ray flux should be $\\sim 0.1, 0.1$, and $0.025$ of the soft X-ray flux in the central 1'' for NGC 1399, NGC 4472, and NGC 4636, respectively.  Deeper {\\it Chandra} observations of NGC 1399 and NGC 4472 may be able to detect emission at this level and would provide interesting constraints on accretion models.  One complication in interpreting such observations is that the black hole mass is estimated from the $M_{bh}-\\sigma$ relation, not direct dynamical studies. This leads to a factor of few uncertainty in $R_A$.  In this respect, M87 is a more promising system since the black hole mass is dynamically determined to be $\\approx 3 \\pm 1 \\times 10^9 M_\\odot$ (e.g., Macchetto et al. 1997); $R_A$ is thus $\\approx$ 1'' as in the Galactic Center.  Unfortunately, however, unresolved emission from the jet appears to dominate {\\it Chandra} observations of the nucleus of M87 (Wilson \\& Yang 2002) and so hard X-ray bremsstrahlung emission will be difficult to detect."
    },
    "0201/astro-ph0201506_arXiv.txt": {
        "abstract": "We present results from  a deep mid--infrared survey of the Hubble Deep Field South (HDF--S) region performed at 7 and 15$\\mu$m with the CAM instrument on board the {\\em Infrared Space Observatory} ({\\em ISO}). The final map in each band was constructed by the coaddition of four independent rasters, registered using bright sources securely detected in all rasters, with the absolute astrometry being defined by a radio source detected at both 7 and 15$\\mu$m. We sought detections of bright sources in a circular region of radius 2.5 arcmin at the centre of each map, in a manner that simulations indicated would produce highly reliable and complete source catalogues using simple selection criteria. Merging source lists in the two bands yielded a catalogue of 35 distinct sources, which we calibrated photometrically using photospheric models of late--type stars detected in our data. We present extragalactic source count results in both bands, and discuss the constraints they impose on models of galaxy evolution models, given the volume of space sampled by this galaxy population.\\\\ ",
        "introduction": "\\begin{table*} \\caption{{\\em ISO} Observation Log. This table gives some details from the {\\em ISO} databases for each of the {\\em ISO} HDF--S observations: The target name, coordinates,  Observation Number (OSN), the time spent on target in seconds (TDT), the revolution number (REV), the status and the date.  Note that two observations (OSN 4 and 8) failed, but were repeated on 27 and 29 November.} \\begin{tabular}{ccccccrr} TARGET       &      RA (J2000)  &  DEC (J2000)   & OSN &  TDT & REV &STATUS     &    Date\\\\~\\\\ \\hline HDF-1 LW2    &    22h 32m 57.5s &-60d 33' 10.0\"  &  1 &  7825 &702 &Observed   &     17 Oct 1997 \\\\ HDF-4 LW2    &    22h 32m 53.9s &-60d 33' 00.0\"  &  7 &  7825 &702 &Observed   &     17 Oct 1997 \\\\ HDF-2 LW2    &    22h 32m 56.4s &-60d 32' 51.8\"  &  3 &  7825 &704 &Observed   &     19 Oct 1997 \\\\ HDF-4 LW3    &    22h 32m 53.9s &-60d 33' 00.0\"  &  8 &  7497 &722 &Failed     &      6 Nov 1997 \\\\ HDF-2 LW3    &    22h 32m 56.4s &-60d 32' 51.8\"  &  4 &  7497 &722 &Failed     &      6 Nov 1997 \\\\ HDF-3 LW2    &    22h 32m 55.0s &-60d 33' 18.2\"  &  5 &  7825 &723 &Observed   &      7 Nov 1997 \\\\ HDF-3 LW3    &    22h 32m 55.0s &-60d 33' 18.2\"  &  6 &  7497 &723 &Observed   &      7 Nov 1997 \\\\ HDF-1 LW3    &    22h 32m 57.5s &-60d 33' 10.0\"  &  2 &  7497 &723 &Observed   &      8 Nov 1997 \\\\ HDF-4 LW3    &    22h 32m 53.9s &-60d 33' 00.0\"  &  8 &  7497 &742 &Observed   &     27 Nov 1997 \\\\ HDF-2 LW3    &    22h 32m 56.4s &-60d 32' 51.8\"  &  4 &  7497 &745 &Observed   &     29 Nov 1997 \\\\  \\end{tabular}  \\end{table*}    One of the most notable achievements of the {\\em Hubble Space Telescope} ({\\em HST}) has been to lead and inspire the concerted multi--wavelength programme of observations of the {\\em Hubble Deep Field} (HDF, Williams et al. 1996)\\nocite{Williams et al. 1996} region. As part of that campaign we observed the HDF at 6.7 and 15 $\\mu$m using the {\\em ISO}--CAM instrument (Cesarsky et al. 1996) on the {\\em Infrared Space Observatory} ({\\em ISO}: Kessler et al. 1996). From the maps that resulted from these observations (Serjeant et al. 1997) we extracted sources in both bands (Goldschmidt et al. 1997), whose number counts implied a strongly--evolving population of starburst galaxies (Oliver et al. 1997).  Following the association of these sources with galaxies in optical HDF catalogues (Mann et al. 1997) we derived an infrared luminosity density that suggested a higher star--formation rate in the HDF region than indicated by optical studies (Rowan-Robinson et al. 1997): the importance of dust obscuration in estimating the star-formation rate has been confirmed by other {\\em ISO} surveys e.g. \\cite{Flores et al. 1999}, from detailed consideration of the optical measures of star formation \\cite{Steidel et al. 1999} and from the intercomparison of different star formation indices \\cite{Cram et al. 1998}   Difficulties with the {\\em ISO} 6.7 $\\mu$m data led us to re--observe the HDF at that wavelength.  These new data, together with a consensus view of the interpretation of our {\\em ISO} HDF--N data derived from the combined experience of the several groups that re--analyzed them (Aussel et al. 1999, Desert et al. 1999) in the light of developing knowledge of the properties {\\em ISO}--CAM data will be the topic a subsequent paper, as will a revised and updated scientific interpretation of the {\\em ISO HDF} data.    Following the success of the  HDF project, a similar programme of {\\em HST} observations was planned for the southern hemisphere, and the region of the Hubble Deep Field South (HDF--S) has become the target of a similarly wide--ranging multi--wavelength programme\\footnote{Details of the HDF--S programme can be found at \\verb+http://www.stsci.edu/ftp/science/hdf/hdfsouth/hdfs.html+} of observations. This paper describes our contribution to that project, through our mapping of the HDF--S with {\\em ISO}--CAM. We mapped the HDF--S WFPC2 fields at both 6.7 and 15 $\\mu$m, as in the northern HDF, but with a slightly different observational strategy (described in Section 2), motivated by our experience with the {\\em ISO HDF} data. Section 3 describes our data reduction procedures, and Sections 4 and 5 source extraction and photometric calibration, respectively. In Section 6 we describe simulations of the data performed to facilitate assessment of the reliability of the source catalogues we present in Section 7 and to compute the effective area of the survey as a function of flux cut, as this is required for computation of the source counts, which is the topic of Section 8. Finally, Section 9 presents a discussion of the results of this paper and the conclusions we draw from them. In an accompanying paper (Mann et al., 2002, hereafter Paper II) we seek associations for these sources in optical/near--infrared and radio surveys of the HDF--S region, and present star formation rate estimates for the sources for which we find associations. ",
        "conclusions": "We have performed a survey using {\\em ISO-CAM} at 6.7 and 15 $\\mu$m in the Hubble Deep Field South region. The observational and data reduction techniques that we have employed mean that these data, and the 6.7$\\mu$m data in particular, are significantly improved over the equivalent {\\em ISO-HDF} data.  From the resulting data we have extracted conservative bright source lists.  We have throughly investigated the completeness and reliability of these lists using simulations.  We have performed an external calibration of the data using stars in the field, a number of which have been spectroscopically classified.  We find that the 6.7 and 15$\\mu$m colour flux diagram provides a useful discriminant between stars and galaxies.  We have investigated the number counts of the extra-galactic sources and stars.  We find that the number counts of the extra-galactic sources are consistent with previous determinations, however, we stress that the volume sampled by our survey is likely to small and so clustering effects (cosmic variance) may mean that this agreement is somewhat coincidental.  A steep upturn at the faintest fluxes is due to only a few sources and is almost certain to be an effect of clustering. Further details of this project can be found at {\\tt astro.ic.ac.uk/hdfs})."
    },
    "0201/astro-ph0201056_arXiv.txt": {
        "abstract": "Old high-$z$ galaxies are important tools for understanding the structure formation problem  and may become the key to determine the ultimate fate of the Universe. In this {\\it letter}, the inferred ages of the three oldest galaxies at high redshifts reported in the literature are used to constrain the first epoch of galaxy formation and to reanalyse the high-z time scale crisis. The lower limits on the formation redshift $z_f$ depends on the quantity of  cold dark matter in the Universe. In particular, if $\\Omega_m \\geq 0.37$ these galaxies are not formed in FRW cosmologies with no dark energy. This result is in line with the Supernovae type Ia measurements which suggest that the bulk of energy in the Universe is repulsive and appears like an unknown form of dark energy component. In a complementar analysis, unlike recent claims favoring the end of the age problem, it is shown that the Einstein-de Sitter model is excluded at high-z by $\\sim 3\\sigma$. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0201/astro-ph0201260_arXiv.txt": {
        "abstract": "Nearby galactic nuclei are observed to be very much dimmer than active galactic nuclei in distant galaxies.  The Chandra X-ray Observatory has provided a definitive explanation for why this is so.  With its excellent angular resolution, Chandra has imaged hot X-ray-emitting gas close to the gravitational capture radius of a handful of supermassive black holes, including Sgr A$^*$ in the nucleus of our own Galaxy.  These observations provide direct and reliable estimates of the Bondi mass accretion rate $\\dot M_{Bondi}$ in these nuclei.  It is found that $\\dot M_{Bondi}$ is significantly below the Eddington mass accretion rate, but this alone does not explain the dimness of the accretion flows.  In all the systems observed so far, the accretion luminosity $L_{acc}\\ll 0.1\\dot M_{Bondi}c^2$, which means that the accretion must occur via a radiatively inefficient mode. This conclusion, which was strongly suspected for many years, is now inescapable.  Furthermore, if the accretion in these nuclei occurs via either a Bondi flow or an advection-dominated accretion flow, the accreting plasma must be two-temperature at small radii, and the central mass must have an event horizon.  Convection, winds and jets may play a role, but observations do not yet permit definite conclusions. ",
        "introduction": "At high redshift, many galactic nuclei are extremely bright, with luminosities in excess of $10^{45}~{\\rm erg\\,s^{-1}}$.  These active galactic nuclei (AGN) are believed to be powered by accretion of gas onto supermassive black holes (SMBHs) at nearly the Eddington rate \\cite{K99}.  The accretion very likely occurs via a Shakura-Sunyaev thin disk \\cite{SS73},\\cite{NT73},\\cite{P81},\\cite{FKR92}.  The Big Blue Bump, which is present in the spectra of all bright AGN, is identified with blackbody emission from the optically thick disk \\cite{KB99}, while the X-ray emission (roughly 10\\% of the luminosity) is thought to be produced by an optically thin corona above the disk \\cite{HM91},\\cite{HM93}.  In addition, most AGN have substantial infrared emission, usually the result of dust reprocessing at a relatively large distance from the SMBH.  Some AGN also have significant radio emission from relativistic jets.  Most galactic nuclei at low redshift are very different.  These nearby nuclei are much less active --- sometimes not active at all.  The nuclear source in our own Galaxy, Sagittarius A$^*$, is a particularly good example of a dim galactic nucleus.  Studies of AGN demographics suggest that a significant fraction of (perhaps all?)  SMBHs must have gone through an AGN ``lighthouse'' phase at some early stage in their lives.  Why and how did these early lighthouses switch off to become the dormant nuclei we see today?  The reduced activity in nearby galactic nuclei is certainly not because of a lack of SMBHs.  Recent observations have provided ample evidence that virtually every galaxy in our local neighborhood has a SMBH \\cite{Retal98}.  The lack of activity must therefore be the result of reduced gas supply.  While this is certainly part of the answer, we shall see that it is not the whole story.  A second, equally important, reason is that the very mode of accretion is different in dim nuclei: the accretion occurs via a radiatively inefficient mode, such that even what little gas reaches the SMBH produces much less radiated energy per unit accreted mass than in bright AGN.  We review below the evidence for this conclusion.  Our primary emphasis is on the nearest and best-studied dim SMBH, Sgr A$^*$, but we also discuss briefly other dim nuclei.  The focus is on the observations and on what we can or cannot tell with confidence from the presently available data. ",
        "conclusions": "The Chandra X-ray Observatory has eliminated a major uncertainty that has hampered our understanding of dim galactic nuclei.  Thanks to Chandra's excellent angular resolution, we now have direct measurements of the density and temperature of ambient gas close to the gravitational capture radius of the SMBHs in Sgr A$^*$ and a few nearby galactic nuclei.  This information allows us to estimate for these nuclei the Bondi mass accretion rate $\\dot M_{Bondi}$, as well as the accretion rate in an advection-dominated accretion flow $\\dot M_{ADAF}$.  With the uncertainty in $\\dot M$ removed, we are now in a position to answer the question posed in the title of this article. The answer consists of three parts:  \\begin{itemize} \\item In all the dim galactic nuclei for which Chandra has provided a direct estimate of $\\dot M_{Bondi}$, the accretion rate is found to be well below the Eddington rate.  This is in contrast to bright AGN which are believed to accrete at close to the Eddington rate.  Thus, the first, and obvious, reason why AGN switch off is that the gas supply to the SMBH is reduced, presumably because most of the gas has been converted into stars. \\item Even after allowing for the reduced $\\dot M$, the objects studied are still anomalously dim: $L_{acc}\\ll0.1\\dot M_{Bondi}c^2$.  Therefore, we can state with great confidence that the accretion {\\it must} proceed via a radiatively inefficient mode.  A two-temperature ADAF model fits the available data quite well (Fig. 4), with almost no adjustable parameters (only $\\delta$ needs to be adjusted, and even it is loosely constrained: \\S5.2).  In this model, there is not much room for additional emission from a jet.  A Bondi model can also be made to fit the data, provided the accreting gas is taken to be two-temperature (\\S4.2).  However, the neglect of angular momentum of the gas is a serious weakness of the model.  Both the ADAF and Bondi models work only if the central object has an event horizon (\\S5.4). Independently of these results, one can state with high confidence that there is no Shakura-Sunyaev thin disk in Sgr A$^*$ (\\S3). \\item Outflows and convection may be important, in which case the mass accretion rate onto the BH may be significantly less than the mass supply on the outside, i.e.  $\\dot M_{BH}\\ll M_{Bondi},\\dot M_{ADAF}$ (see the discussion of ADIOS/CDAF models in \\S\\S7.1--7.3).  There is as yet no compelling observational evidence for these models, but there are strong theoretical reasons for favoring them.  If the accretion flows in Sgr A$^*$ and other dim nuclei are of the ADIOS or CDAF type, then the accretion flow may be very dim in radio/mm, and the observed emission in these bands may come from a relativistic jet or some other component external to the accretion flow. \\end{itemize}  Thus, it appears that three different effects all conspire to make nearby galactic nuclei extraodinarily dim: there is less gas available, the gas accretes via a radiatively inefficient mode, and (perhaps) less gas reaches the BH than is available for accretion.  \\bigskip\\noindent \\emph{Acknowledgements:} The author thanks Shin Mineshige and Eliot Quataert for useful comments on the manuscript and the W.M. Keck Foundation for support as a Keck Visiting Professor at the Institute for Advanced Study, Princeton.  This research was supported by NSF grant AST-9820686."
    },
    "0201/astro-ph0201326_arXiv.txt": {
        "abstract": "We report on optical and near-infrared (NIR) follow-up spectroscopy of faint far-infrared (FIR) sources found in our deep FIR survey by Kawara et al. ",
        "introduction": "Deep surveys at FIR and submilimeter wavelengths have been carried out in order to investigate the nature of dust-enshrouded galaxies at high redshift. As a contribution to this field, our group made a deep FIR survey using the ISOPHOT camera on board the {\\it Infrared Space Observatory} (ISO) satellite (Kawara et al. 1998; Matsuhara et al. 2000). Mapping at 90$\\mu$m and 170$\\mu$m of two $44^\\prime \\times 44^\\prime$ fields in the Lockman Hole (LH\\_EX and LH\\_NW), a region exhibiting the lowest H{\\ts}{\\sc i} column density in the sky (Lockman et al. 1986), resulted in the detection of 36 sources with $f_{90} >${\\ts}150{\\ts}mJy and 45 sources with $f_{170} >${\\ts}150{\\ts}mJy.  Given the relatively large size of the ISOPHOT beam at 170$\\mu$m ($\\sim$90$^{\\prime\\prime}$), we have obtained opt/NIR images and spectra using telescopes on Mauna Kea and 6cm radio continuum maps using the VLA (Yun et al. 2002) to identify the most likely source of the 170$\\mu$m emission.  Here we report our initial identifications of the brightest of the ISOPHOT 170$\\mu$m sources. ",
        "conclusions": "Redshifts of 35 FIR source candidates were determined using optical spectra obtained with ESI on Keck{\\ts}II during three observing runs in 2000 March and 2001 January.  Infrared luminosities, $L_{\\rm ir}(8-1000\\mu m)$, were then estimated by using the ISOPHOT fluxes and assuming an SED similar to that of Arp{\\ts}220.  We found one hyperluminous infrared galaxy (HyLIG: $L_{\\rm ir} > 10^{13} L_\\odot$) at $z=1.6$, 11 ultraluminous infrared galaxies (ULIGs: $L_{\\rm ir} > 10^{12} L_\\odot$) at $0.3 < z < 1$, 12 luminous infrared galaxies (LIGs: $L_{\\rm ir} > 10^{11} L_\\odot$), and 11 galaxies with $L_{\\rm ir} < 10^{11} L_\\odot$. Except for one LIG at $z=0.365$, all of the galaxies with $L_{\\rm ir} < 10^{12} L_\\odot$) are at $z<0.3$.  The mean redshift for all sources is $0.31\\pm0.31$.  The low-resolution ESI spectra were used to determine the optical spectral-type of the candidate ISOPHOT sources. Following procedures used by Murayama \\& Taniguchi (1998), the spectra were classified into four types -- AGNs, LINERs, HII-type, and early-type (without emission lines).  The HyLIG at $z=1.6$ was found to be a quasar.  One ULIG had an early-type spectrum and 10 ULIGs are HII galaxies.  Among the remaining 23 lower-luminosity sources, there was one early-type galaxy, one Seyfert 2, 10 LINERs and 11 HII galaxies. Thus, based on our low-resolution ESI optical spectra most of the ISOPHOT 175$\\mu$m sources appear to be powered primarily by star formation, consistent with the conclusion reached from an analysis of ISOPHOT number counts by Matsuhara et al. (2000) that most of the ISOPHOT sources are star-forming galaxies at $z<1$.  \\begin{figure} \\plotfiddle{fig1.ps}{250pt}{0}{50}{50}{-150}{-68} \\caption{Low-resolution emission line diagostics of ISOPHOT source candidates.} \\end{figure}  \\vskip 0.1cm"
    },
    "0201/astro-ph0201332_arXiv.txt": {
        "abstract": "It is shown that the low-mass groups obey the $L_{x}\\sim\\sigma_v^{4}$ law deduced for galaxy clusters. The impression of the more shallow slope of the $L_x-\\sigma_{v}$ correlation for groups is created not by enhanced X-ray emission, but by underestimation of the radial velocity dispersion of some groups. ",
        "introduction": "Solinger \\& Tucker (1972) showed that if the source of the X-ray radiation is hot gas bound in clusters, then the X-ray luminosity, $L_{x}$, should be correlated with the radial velocity dispersion, $\\sigma_v$. Thermal emission from the intracluster gas yields an X-ray luminosity, $L_{x}$, proportional to the square of the gas density. In the case of a constant mass-to-light ratio, the $L_{x}$ is proportional to the square of the mass of the cluster. If the cluster is a relaxed system, $\\sigma_v$ is roughly proportional to the square root of the mass. Thus, $L_{x} \\propto \\sigma_v^{4}$. Quintana \\& Melnick (1982) showed that $L_{x}$ of galaxy clusters, indeed, obeys the expected correlation. Dell`Antonio et al. (1994) showed that rich groups follow the $L_x \\propto \\sigma_{v}^4$ relation, but groups with smaller $\\sigma_v$ do have more shallow slope, $L_{x} \\propto \\sigma_{v}^{2.7}$. All recent observations (Mahdavi et al. 1997, Zabludoff \\& Mulchaey 1998, Helsdon \\& Ponman 2000, Mahdavi et al. 2000, Xue \\& Wu 2000, Mahdavi et al. 1997) proved that the dependence of the X-ray emission for low-mass groups (which includes compact groups) on the $\\sigma_v$ is much weaker than for galaxy clusters. Zimer et al. (2001) analyzing the results of different investigators mentioned some discrepancies, and concluded that they were due to poorly determined $\\sigma_v$s and $L_{x}$s. All data, show, however, that some amount of low-mass groups of galaxies are located on the left side of the line $L_x \\propto \\sigma_{v}^{\\sim4}$. It has been generally assumed that the reason of such location of groups on the graph $L_x - \\sigma_{v}$ is the enhanced (by one-two orders of magnitude) X-ray emission of groups. It is widely assumed that the excess X-ray luminosity of the low-mass groups is explained by the \"mixed emission\" scenario (Dell`Antonio et al. 1994) when the emission from the intragroup plasma may be contaminated by a superposition of diffuse X-ray sources corresponding to the hot interstellar medium of the member galaxies. However, the shift of groups to the left of the $L_x \\propto \\sigma_{v}^{4}$ line may have another reason: underestimation of $\\sigma_v$s.  It has been shown that HCGs and ShCGs have a triaxial spheroid, \"cigar\"-like shapes (Malykh \\& Orlov 1986, Hickson et al. 1984, Oleak et al. 1998). In a series of papers Tovmassian and collaborators (Tovmassian et al. 1999, Tovmassian \\& Chavushyan 2000, Tovmassian et al. 2001) showed that $\\sigma_v$s of compact groups (CGs), and of associated with them loose groups (LGs) which are also elongated and have the same orientation as corresponding CGs, are correlated with elongation of groups determined by $b/a$ ratio\\footnote{$a$ is the angular distance between the most widely separated galaxies in the group, and $b$ is the sum of the angular distances $b_{1}$ and $b_{2}$ of the most distant galaxies on either side of the line $a$ joining the most separated galaxies (Rood 1979).}. It means that members of CGs and LGs move along the elongation of corresponding group. It has been shown by Tovmassian (2001a, 2001b) and Tovmassian \\& Tiersch (2001) that out of three possibilities of such movement: flying out of galaxies from the center of the group in opposite directions, infalling from opposite directions, and regular rotation of member galaxies in elongated orbits around the gravitational center of each system, the latter possibility is the more realistic one. The rotation time is less than $\\sim3\\times10^9$ years, so CG+LG systems may well be virialized. The measured $\\sigma_{v}$s of such elongated groups depend on the orientation of the group. The highest values of $\\sigma_{v}$s are observed in those groups orientation of elongation of which is close to the line of sight. Such groups generally have the highest $b/a$ ratio, though the chain-like groups oriented at small angles $\\theta$ to the line of sight would also have relatively high $\\sigma_{v}$. The measured $\\sigma_{v}$s of the majority of groups oriented at intermedient angles to the line of sight are smaller. The groups oriented close to the orthogonal to the line of sight have the smallest measured $\\sigma_{v}$s, i.e. their $\\sigma_{v}$s are highly underestimated. The latter groups are the most elongated and have the smallest $b/a$ ratio. Hence, they would be located on the left part of the $L_{x}-\\sigma_{v}$ graph. We show in this paper that underestimation of $\\sigma_{v}$s of the seen edge-on elongated groups, indeed, creates the more shallow slope of the $L_{x}-\\sigma_{v}$ correlation. ",
        "conclusions": "Consideration of the $b/a$ ratios of the low-mass groups shows that groups located at the utmost left of the $L_x \\propto \\sigma_{v}^4$ line on the $L_x - \\sigma_{v}$ graph have, on average, smaller $b/a$ values than those located at the utmost right. The elongation of groups with small $b/a$ ratios are oriented close to the orthogonal to the line of sight. The measured $\\sigma_{v}$s of these groups are, thus, underestimated, and are smaller than the real values. Therefore, such groups are artificially shifted to the left on the $L_x - \\sigma_{v}$ graph. This creates an impression of an enhanced X-ray luminosity. If to take into account the reasonable amount of underestimation of $\\sigma_{v}$s (of the order of 200-300 km s$^{-1}$, which corresponds to angle $\\theta$ of about $40\\deg-50\\deg$), then the corresponding groups will be moved to the right, towards the line $L_x \\propto \\sigma_{v}^4$. Hence, the low-mass groups {\\it obey} the $L_x \\propto \\sigma_{v}^4$ law for clusters of galaxies (Solinger \\& Tucker 1972, Quintana \\& Melnick 1982). It means that there is no need to apply any mechanism of the enhancement of the X-ray luminosity of the low-mass groups, since in reality there is no any enhancement.  Compact groups are stable systems with members probably rotating around the gravitational center of the corresponding group (Tovmassian 2001a, 2001b). For the reason of regular movement in elongated orbits the velocities of member galaxies in the central region of a CG are high enough. Therefore, the efficiency of interaction in such groups would be smaller than in the made numerical simulations when such regular movement has been neglected (Barnes 1985, 1989; Ishizawa 1986; Mamon 1987, 1990; Zheng et al. 1993). The formation of the hot interstellar medium in member galaxies, widely assumed for explanation of the excess X-ray emission, may be very rare, and may not dominate the global X-ray emission. In fact, the claimed by many excess of the X-ray luminosity in the low-mass groups is due to a projection effect, and is, thus, a result of {\\it misinterpretation} of the observational data. The finding of Mahdavi \\& Geller (2001) that clusters of galaxies and single elliptical galaxies form a continuous relation $L_x - \\sigma_{v}^m$ are consistent with the result presented in this paper. Low-mass groups are not exotic objects and obey the same law."
    },
    "0201/astro-ph0201104_arXiv.txt": {
        "abstract": "We have used the IRAM Plateau de Bure mm interferometer to locate with subarcsecond accuracy the dust emission of three of the brightest 1.2mm sources in the NTT Deep Field (NDF) selected from our 1.2mm MAMBO survey at the IRAM 30m telescope. We combine these results with deep B to K imaging and VLA interferometry.  Reliable identifications are an essential step towards an understanding of the high redshift (sub)mm galaxy population, towards testing the common belief that they are scaled up analogs of local dusty ultraluminous galaxies, and in shedding light on the possible connection to spheroid formation.  Strikingly, none of the three accurately located mm galaxies MMJ120546-0741.5, MMJ120539-0745.4, and MMJ120517-0743.1 has a K-band counterpart down to the faint limit of K$_s$$>$21.9.  This implies that these three galaxies are either extremely obscured and/or are at very high redshifts (z$\\ga$4). We combine our results with literature data for 11 more (sub)mm galaxies that are identified with similar reliability. In terms of their K-band properties, the sample divides into three roughly equal groups: (i) undetected to K$\\sim$22, (ii) detected in the near-infrared but not the optical and (iii) detected in the optical with the possibility of optical follow-up spectroscopy. We find a trend in this sample between near-infrared to submm and submm to radio spectral indices, which in comparison to spectral energy distributions (SEDs) of low redshift infrared luminous galaxies suggests that the most plausible primary factor causing the extreme near-infrared faintness of our objects is their high redshift. We show that the near-infrared to radio SEDs of the sample are inconsistent with SEDs that resemble local far-infrared cool galaxies with moderate luminosities, which were proposed in some models of the submm sky. We briefly discuss the implications of the results for our understanding of galaxy formation. ",
        "introduction": "A local census of the distribution of the baryonic mass reveals, albeit with large uncertainties, that of the baryons presently locked in stars, a majority reside in spheroids \\citep{per92,fuk98}. A key question then is when and how all these baryons have come to reside in spheroids. For several decades there have been two competing explanations. The classical pictures are the 'monolithic collapse' of \\citet{egg62} versus the (hierarchical) merging model of \\citet{sea78}.  With the development of new techniques \\citep[e.g., galaxies which 'drop-out' in deep images;][]{ste96} and new technology (e.g., 10m class telescopes and imaging mm/submm bolometers) our understanding of the cosmic history of star-formation has accumulated rapidly and we are now in a position to begin to address the question as to how and when spheroids were assembled.  A new route for the investigation of star-formation at the highest redshifts (where spheroids probably have been assembled) has been opened by surveys of the dust reradiation in the submm/mm. Recently such surveys have been carried out with the (sub)mm array cameras SCUBA and MAMBO \\citep[e.g., ][]{sma97,hug98,bar98,eales99,ber00a}. These investigations detect a population of luminous high redshift infrared galaxies that represent a significant fraction of the cosmic submm background detected by COBE. A unique advantage is the near constancy of observed (sub)mm flux with redshift because of the favorable negative k-correction for galaxy SEDs peaking in the rest frame far-infrared. As such, (sub)mm observations favor the detection of objects resembling extreme versions of the most luminous infrared sources in the local universe, ultraluminous infrared galaxies (ULIRGs) -- either powerful heavily obscured star-formation ($\\sim$ 10$^2$ -- 10$^3$ M$_{\\sun}$ yr$^{-1}$) or AGN.  Given their likely extreme luminosities and the fact that a majority of the submm sources do not appear to be AGN \\citep[e.g.,][]{fab00,bau00,hor00,bar01a,bar01b,alm01}, it is hard to resist the speculation that a substantial fraction of faint (sub)mm population are massive spheroids in formation \\citep[e.g.,][]{san99,fra99,lil99,tan99}. Analysis of the stellar content of spheroids \\citep[e.g.,][]{tra00a,tra00b} and investigations of the dynamics of merging galaxies hosting powerful starbursts at low redshift \\citep{gen01} suggest diverse formation mechanisms and evolutionary histories of early type galaxies.  If the speculation of the connection between spheroid formation and the faint submm population turns out to be well-founded, then determining the redshifts, bolometric luminosities, and morphological and spectral properties of a significant number of sources would greatly enhance our understanding of how spheroids formed and which mechanism (monolithic collapse or merging) most strongly influenced how and when spheroids formed.  Beginning in the winter of 1998, our groups at the MPIfR, MPE, and NRAO have been conducting a deep (rms $\\sim$ 0.5mJy), wide (each of three fields more than 100 arcmin$^{2}$) area survey at 1.2mm with the Max-Planck-Millimeter Bolometer Array \\citep[``MAMBO'';][]{kre98} on the IRAM 30m telescope.  These surveys are specifically designed to detect significant numbers of the brightest mJy-sources at 1.2mm \\citep{car99a,car00b,ber00a,ber00b}.  The large areal coverage and high mapping speed of MAMBO have produced more than fifty firm detections of bright mm sources (Bertoldi et al. 2002, in preparation), which directly implies high luminosity through the strong negative k-correction for sources that have spectral energy distributions peaking in the infrared.  One region that we have surveyed with MAMBO is a southern field centered on, but larger than, the NTT Deep Field \\citep[NDF;][]{arn99}.  Here we report on mm interferometry, radio interferometry, and deep K-band and BVRI imaging of three of the brightest likely nonlensed MAMBO mm sources in the NDF (S$_{1.2mm}>$3mJy; for the perhaps more familiar S$_{850\\mu m}$ scale this converts to $\\ga$8mJy for z$\\sim$3 but depends on SED and redshift). Our strategy is to identify possible counterparts and to begin to assemble spectral energy distributions constraining the nature and redshift of these objects. The interferometric observations with the IRAM Plateau de Bure Interferometer (PdBI) and the VLA are required to determine positions of subarcsecond accuracy and fluxes. Several faint potential optical/near-IR counterparts are often present inside or near the MAMBO beam for a mm galaxy \\citep[e.g.,][]{lil99}. An accurate location by interferometry is essential for any identification because of this ambiguity and the possibility of `blank fields'. Without that step, even the existence of potential counterparts does not imply that the mm source is in fact associated with one of them.  We discuss our results in the context of the data for other reliably identified sub(mm) galaxies.  We adopt the cosmological parameters: $\\Omega_{matter}=0.3$, $\\Omega_{\\Lambda}=0.7$, and H$_0$=70 km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "We have used mm and radio interferometry to locate accurately three of the brightest and likely non-lensed MAMBO mm sources -- MMJ120546-0741.5, MMJ120539-0745.4, and MMJ120517-0743.1.  These objects are among the very brightest (sub)mm galaxies if amplification is taken into account. Associated deep optical and K-band follow-up shows that these sources do not have optical or near-infrared counterparts down to very faint magnitudes (K$_{s}$$>$21.9).  The near-infrared faintness of these bright mm galaxies implies that they must either be at very high redshift (z$\\ga$4) or more highly obscured than even the dustiest local ULIRGs.  We have combined these limits with results from the literature for those among the brightest (sub)mm sources that we consider located reliably to subarcsecond accuracy. About a third each of the total sample are (1) detected in the optical with follow-up possibility, (2) optically faint but detected at K$\\sim$20-21, and (3) undetected in K at levels similar to the three MAMBO galaxies (K$\\ga$22). The combined optical -- near-infrared -- submm -- radio data for the sample suggest the near-infrared faintness of our objects to be most likely due to very high redshifts and not due to obscuration alone.  The SEDs of local galaxies that are cool in the far-infrared do not explain the overall characteristics of these galaxies very well and therefore models of this type are not favored on the basis of these data.  The nondetection of a major fraction of the (sub)mm population in the near-infrared and optical stresses the need for future new methods like wideband CO searches, mid-infrared emission features, or maser lines in obtaining accurate redshifts for this important group of high redshift objects, in order to understand the (sub)mm galaxy population as a whole.  The existence of a sizeable fraction of K-band faint (sub)mm galaxies has implications for our understanding of the formation of galaxies. Specifically, such results will test whether the faint mm population is similar to the merger-induced dusty powerful starbursts in the local universe, or if they represent something more extreme -- e.g., the ``monolithic collapse'' of massive spheroids.  If the brightest mm sources lie mostly at high redshift (z$\\ga$4) and have extreme luminosities (L$_{bol}$$>$ few $\\times$ 10$^{12}$ L$_{\\sun}$), then hierarchical merging models will have difficulty in explaining these sources within the general context of galaxy formation and evolution.  As such, these bright source may provide a critical test of these widely favored models."
    },
    "0201/astro-ph0201274_arXiv.txt": {
        "abstract": "Promising methods for studying galaxy evolution rely on optical emission line width measurements to compare intermediate-redshift objects to galaxies with equivalent masses at the present epoch. However, emission lines can be misleading. We show empirical examples of galaxies with concentrated central star formation from a survey of galaxies in pairs; HI observations of these galaxies indicate that the optical line emission fails to sample their full gravitational potentials. We use simple models of bulge-forming bursts of star formation to demonstrate that compact optical morphologies and small half-light radii can accompany these anomalously narrow emission lines; thus late-type bulges forming on rapid (0.5~--~1 Gyr) timescales at intermediate redshift would exhibit properties similar to those of heavily bursting dwarfs. We conclude that some of the luminous compact objects observed at intermediate and high redshift may be starbursts in the centers of massive galaxies and/or bulges in formation. ",
        "introduction": "Optical emission line widths are potentially important diagnostic tools for measuring the intrinsic gravitational masses of galaxies within their optical radii.  Because of the sensitivity requirements for spatially resolved rotation curves, large surveys of galaxies at intermediate redshift and studies of galaxies at high redshift use unresolved or ``integrated'' emission line widths, computed from Gaussian fits to emission lines in the spectrum of the whole galaxy.  However, the results are sometimes ambiguous in the case of compact star-forming galaxies.  The luminous compact blue galaxies observed at intermediate redshift (e.g., Koo et al. 1994; 1995) have small half-light radii (R$_{\\rm e}=1$~--~3.5 kpc) and narrow emission-line velocity widths ($35 < \\sigma < 126$~km~s$^{-1}$). These properties suggest that although they are luminous galaxies, compact blue galaxies may be intrinsically faint galaxies undergoing a strong burst of star formation (Guzm\\'{a}n et al. 1996; 1997).  However, Kobulnicky \\& Zaritsky (1999) measure high metallicities for these objects, appropriate only for massive galaxies, and HST images show possible evidence for surrounding older populations (Guzm\\'{a}n et al. 1998). Similarly, at higher redshifts ($z \\sim 3$) the ``Lyman break'' galaxies also exhibit narrow integrated line widths that do not correlate with galaxy luminosity (Pettini et al. 2001). These observations taken together raise the question of whether emission line widths of compact objects accurately trace their potential wells (Kobulnicky \\& Zaritsky 1999).  We present evidence from observations of local galaxies and simple models of compact star formation that centrally concentrated star formation changes the measured emission line widths and half-light radii of galaxies (see also Kobulnicky \\& Gebhardt 2000; Pisano et al. 2001). This star formation can arise from major mergers (Mihos \\& Hernquist 1996), minor mergers (Mihos \\& Hernquist 1994), and secular evolution (Pfenniger \\& Norman 1990), processes that may be directly linked to bulge formation.  The number of compact blue objects at intermediate redshift that are actually concentrations of star formation in larger galaxies remains unknown. If the luminous, compact blue galaxies are frequently bulges in formation, their number counts contain information about the timescales for evolution along the Hubble sequence. ",
        "conclusions": ""
    },
    "0201/astro-ph0201042_arXiv.txt": {
        "abstract": "We have analysed images of the field of A2390 obtained with the CFHT and HST. The analysis fits models to bulge and disk components to several hundred galaxies, with about equal samples from the cluster and field. We also have assessed and graded asymmetries in the images. The cluster galaxies are compared in different cluster locations and also compared with field galaxies. We find that the central old population galaxies are bulge-dominated, while disk systems have young populations and are found predominantly in the outer cluster. S0 and bulgy disk galaxies are found throughout, but concentrate in regions of substructure. Disks of cluster blue galaxies are generally brighter and smaller than those in the field. We find that the cluster members have a higher proportion of interacting galaxies than the field sample. Interactions in the cluster and in the field, as well as cluster infall, appear to inhibit star-formation in galaxies. ",
        "introduction": "The well-known rich cluster Abell 2390 at z=0.23 was studied as part of the CNOC1 cosmology program (Carlberg et al 1996). A detailed discussion of its spectroscopic properties from the CNOC1 database was published by Abraham et al (1996). From spectra of some 250 cluster galaxies, Abraham et al showed that infall into the cluster truncates star-formation, although there are some `near-field' galaxies that have young populations, and may not yet be affected by the cluster environment.  The morphological structure of cluster galaxies is of interest in further understanding the effects of cluster infall and assimilation of galaxies. It is presumed that the high fraction of red central galaxies with old stellar populations are formed by tidal events that occur within the cluster, and giving them a bulge-dominated morphology. However, the CNOC1 imaging data were not used for a systematic morphological study until now.  In this paper, we have made use of a variety of additional imaging data of the cluster. We have performed a morphological modelling analysis of the cluster galaxies, and also of a field galaxy sample from the same databases. In our discussion we discuss the connection between morphology, cluster location, and spectra, and look for differences of properties in a similar sample of field galaxies.  The data used are from four databases. 1) the original CFHT MOS images from the CNOC1 database. These cover all the galaxies which have spectra, so are complete. However, the image quality and sampling are not optimal. 2) CFHT OSIS images have better sampling and image quality, from better observing conditions and with a tip-tilt imaging system. However, these images were part of a narrow-band imaging program and were not deeply exposed in the continuum band. 3) Images with the CFHT AO system in visible wavelengths produced still better sampled images with FWHM in the range 0.3 to 0.4\". 4) Finally, we used HST WFPC2 images from our own and archival databases. These have FWHM about 0.2-0.3\".  Table 1 summarizes the data and compares their coverage of the cluster. While the MOS data have the worst resolution, they offer the most extensive coverage of the cluster and have the highest signal level per pixel. Thus, in what follows, we have used the MOS data for most of the discussion, but have used the other datasets as checks on the morphology results. The analysis does create a PSF that may vary across the image fields, and uses them to calibrate the images, so that image resolution is not the overriding consideration for the images. While the OSIS data have superior resolution and sampling, their coverage of the cluster is less complete and the signal levels are much lower, so that they are not as useful. The PUEO and WFPC2 data are good quality checks on a limited subsample of galaxies. ",
        "conclusions": ""
    },
    "0201/nucl-th0201026_arXiv.txt": {
        "abstract": "We propose a new nucleosynthesis process in which various nuclei are formed through the liquid-gas phase transition of supernova matter as a preprocess of the standard r-process, and study its possibility of realization qualitatively. In the relativistic mean field and statistical model calculations, we find that this process may take place and help the later r-process to proceed. ",
        "introduction": "It is generally believed that there exist several phases of nuclear matter. Among the phase transitions between them, the nuclear liquid-gas phase transition has been extensively studied in these three decades. It takes place at relatively cold ($T_c = (5-8)$ MeV) and thin ($\\rho_B \\sim \\rho_0/3$) nuclear matter, and it causes multifragmetation of expanding nuclear matter in heavy-ion collisions~\\cite{HeavyIon}. In the universe, the temperature and density region of this transition would be probed during supernova explosion. Since supernova explosion can be regarded as expansion of infinite matter composed of nucleons and leptons, its fragmentation through the phase transition may produce various nuclei including heavy nuclei which are believed to be synthesized in r-process or rp-process. Then the phase transition during supernova explosion may be helpful to solve some problems in these standard processes.  In this work, we propose a new kind of nuclear synthesis process through the liquid-gas phase transition (LG process) as a preprocess of the usual r-process. In the early stage of explosion, since the density and the temperature are high enough to keep statistical equilibrium and to trap neutrinos, entropy per baryon $S/B$ and lepton to baryon ratio $Y_L$ are to be conserved. This condition is expected to be satisfied within the neutrino sphere ($\\rho_B=(10^{-4}-1)\\rho_0$). At these densities, supernova matter expands almost adiabatically and cools down. If the critical temperature $T_c$ of the liquid-gas phase transition is higher than the freeze-out temperature, this matter will fragment and form various nuclei in a critical manner through the phase transition. As the baryon density decreases, charged particle reactions are hindered then the chemical equilibirum ceases to be kept, namely the system freezes out at this point. The statistical distribution of fragments at freeze out will give the initial condition of the later processes such as the r-process.  Although the importance of the liquid-gas phase transition in supernova matter was already noticed and extensively studied before~\\cite{Lattimer}, the main interest in these works was limited to the modification of the equation of state (EoS), then the nuclear distribution as an initial condition of the r-process was not studied extensively. In addition, the mean field treatment assuming one dominant fragment configuration was applied rather than to consider the statistical ensemble of various fragment configurations, which is important at around the critical temperature for sub-saturation densities.  In this study, we first investigate the liquid-gas coexistence region in the $(\\rho_B, T)$ diagram and the ejection possibility of materials inside this coexistence region within the Relativistic Mean Field (RMF) theory. Next, by using a statistical model, we evaluate the fragment formation at around the phase transition and in the coexistence region. Then we compare the fragment distributions at various $(\\rho_B,T)$ points with the solar abundance. ",
        "conclusions": "In this paper, we have proposed the LG process which may play an important role for the understanding of mass and isotope distribution in the unverse. In order to evaluate its possibility qualitatively, we have demonstrated that the coexistence phase extends to low densities keeping the critical temperatures to be greater than 1 MeV, then materials in the coexistence region can be ejected to the outer space at supernova within a two phase model of RMF. Next, we have performed a statistical model calculation which includes the mass table extension and the Coulomb energy correction to adapt to supernova matter. The calculated results show that nuclei above the iron peak would be formed in the coexistence region at relatively high freeze-out densities, and at these densities, observed isotope distribution are explained well.   Works for a more quantitative discussion are in progress."
    },
    "0201/astro-ph0201512_arXiv.txt": {
        "abstract": "We employ an effective gravitational stellar final collapse model which contains the relevant physics involved in this complex phenomena: spherical radical infall in the Schwarzschild metric of the homogeneous core of an advanced star, giant magnetic dipole moment, magnetohydrodynamic material response and realistic equations of state (EOS). The electromagnetic pulse is computed both for medium size cores undergoing hydrodynamic bounce and large size cores undergoing black hole formation. We clearly show that there must exist two classes of neutron stars, separated by maximum allowable masses: those that collapsed as solitary stars (dynamical mass limit) and those that collapsed in binary systems allowing mass accretion (static neutron star mass). Our results show that the electromagnetic pulse spectrum associated with black hole formation is a universal signature, independent of the nuclear EOS. Our results also predict that there must exist black holes whose masses are less than the static neutron star stability limit. ",
        "introduction": "What is the measurable signal for gravitational collapse? It certainly produces gravitational waves, but unfortunately these have shown to be undetectable so far. However, there is another signal. Since the star has a sizable magnetic field, and since we expect this field to get quenched very rapidly during gravitational collapse, a massive electromagnetic pulse (MEMP) is created. We further show that the power spectrum is unlike any other naturally occurring phenomena: a wave packet whose wavetrain lasts only 2 milliseconds and whose characteristic power spectrum is a square block.  We present a simulation, within an effective model, of the iron-core collapse of massive stars which produce type II supernova. Depending on the core's mass, either a hydrodynamic bounce occurs above nuclear densities or the collapse races unhindered into the formulation of a stellar size black hole. Either case produces a MEMP. Our main contribution in this paper is the calculation of the energy spectrum of the MEMP. In cases where the collapsing core has mass greater than the dynamical neutron star mass limit (discussed herein), a hydrodynamic bounce cannot occur and no visible supernova results. In such cases, the MEMP is the only electromagnetic signal associated with the gravitational collapse.  One of the unexpected secondary results is the finding that, in general, there exist two classes of neutron stars: the first class is made up of solitary neutron stars which do not have the opportunity to accrete matter from a binary companion, and therefore have original masses coming from the collapse process. This mass is bounded by the ability of the nuclear equation of state to produce a hydrodynamic bounce. An unexpected secondary finding is that this imposes an heretofore unappreciated constraint on candidate nuclear equations of state. The maximum mass of the neutron star in this class is the dynamical mass, being the largest core mass that can undergo a hydrodynamic bounce. We will see that this dynamical mass limit is considerable less than the static mass limit, so the bounce does indeed save the core from further collapse due to its own self-gravity. The second class of neutron stars are those that can accrete matter from a binary companion and so their maximum mass is constrained by the requirement of static stability: the maximum static neutron star mass. In general, we find that the maximum static mass can exceed the dynamic mass by as large as 100\\%.  A current review of the status of gravitational collapse is given in Joshi (2000), who discusses spherically symmetric collapse. In the present paper, we consider spherical radial infall. As shown in the work of Bocquet et al. (1995), only gigantic magnetic fields B $> \\; 10^{10}$ T = 10 GT (testla = $10^{4}$ gauss, GT = $10^{9}$T) cause stellar deformation. Since the magnetic fields considered in this paper fall far below this large critical field, the star is not deformed by its magnetic field. An interesting followup to the research presented here, is to apply the effective gravitational collapse model to stars having non-negligible rotation.  In general, there are three sources for electromagnetic and gravitational radiation associated with stellar collapse: the direct radiation phase emitted by the stellar object before the formation of a black hole, the so-called damped oscillations known as quasi-normal ringing (Iyer 1987) that are the vibrations of a black hole shaking off its non-zero multipole charge moments, and the late-time power-law tail (Price 1972; Cunningham 1978; Leaver 1986; Hod 1999, 2000) underneath the damped oscillations. Because of the gravitational redshift, essentially the only energy that survives to infinity comes from the direct radiation phase, before the black hole forms. In this paper, we calculate the direct phase electromagnetic energy radiated by stellar objects that bounce and become stable neutron stars, and stellar objects so massive that they become black holes.  In order to calculate the electromagnetic radiation produced by the gravitational collapse, our strategy is the following. During the collapse the infall kinetic energy decreases while the system does mechanical work against the internal star pressure. This pressure has, in general, two components: material (nuclear) pressure and electromagnetic pressure. However, all published nuclear equations of state have the property that the nuclear pressure is orders of magnitude larger than the electromagnetic pressure; indeed, only for magnetic fields greater than $10^{18}$ Gauss does the electromagnetic pressure approach parity with the nuclear pressure. Using conservation of energy, the continuity equation and the material equation of state, we can calculate how the star radius changes with time. If we further assume that the stellar magnetic field can be approximated by a dipole field, and that the lines of magnetic field are frozen in the material and carried along with it, we can relate the change in the star's radius with time to the change in its effective magnetic moment. The usual dipole radiation formula then gives the energy spectrum of the radiation. ",
        "conclusions": "We have presented an effective gravitational collapse model that is thought to include the essential physics: spherical radial infall in the Schwarzschild metric of a homogeneous core of an advanced star possessing a giant magnetic dipole moment. The electrodynamic equations are represented by the approximation of magnetohydrodynamic material response, which will be an excellent approximation due to the high electrical conductivity. Any nuclear EOS can be used, and we chose two which have high interest: the \"Argonne\" AV14+UVII which represents the typical phenomenological nuclear EOS and the QCD EOS. Surprisingly, the maximum core mass that bounced (the dynamic neutron star limit) is rather small using the phenomenological EOS.  The main result of the paper is the calculation of the MEMP from the collapse. Both the total energy and its spectral characteristics are derived. In cases where the core collapse is greater than the dynamical neutron star mass, the MEMP is the only electromagnetic experimental signature. It is hoped that this will spur activity to develop a suitable experimental receiver\\footnote{Due to the peak of the spectral curve around the wavelength $\\sim $ 2km, the receiver will have to be a satellite able to detect a broadband spectrum within the electromagnetic pulse. The additional benefit is that terrestrial signals in this frequency regime cannot penetrate earth's ionosphere, thereby reducing spurious man-made noise.}. To answer the question as to whether the surrounding material from the collapsing object will quench or change an observational signal is clearly complicated, and it requires a detailed calculation in transport theory using information on the supernova models that vary tremendously on their precursor environment, which is beyond the scope of the present paper.  The existence of two classes of neutron star masses leads to the situation where a stellar core greater than the dynamical mass, but less than the static mass, collapses to form a black hole, thus producing black holes less massive than the mass limit of neutron star stability. For example, a black hole of 1.6 solar masses can exist. This possibility sheds light on one of the most perplexing problems of astrophysics: How do black holes of stellar masses form? As envisioned here, a core more massive than the dynamical mass collapses and forms a black hole. It may now accrete the remaining outlying surface matter to produce a whole continuum of stellar mass size black holes.  After this paper was submitted for publication, we were informed of the work of Hanami \\cite{Hanami}. This author attempted to explain gamma-ray bursts using the change in the magnetic field of a collapsing stellar object. Unfortunately, there are several problems associated with this work. Hanami did not consider the fact that the star radiates continuously throughout the collapse trajectory and his numerical solution violates Maxwell's equations. He used a zero pressure EOS, which will not give the correct surface acceleration and core remnant. In order to compute the total energy, the gravitational redshift must be properly included. The calculation of the gravitational redshift is a very difficult problem as explained in the present paper, because, as the star collapses and radiates, each instantaneous stellar configuration has a different redshift. In order to do this, one needs to employ multiple Roll-on-Roll-off mathematical functions in the complex energy plane, correct for the gravitational redshift and add the spectrums incoherently. Hanami did not consider this and simply calculated the redshifted power at the initial collapse configuration, to obtain his result. At the bottom of page 688 and top of page 689, Hanami \\cite{Hanami} alludes to the fact that his solution is the quasi-normal ringing (QNR), but his eigenfrequency is incorrect (QNR have precise discrete eigenfrequencies). Also his time dependence does not correspond to any known quasi-normal mode damping."
    },
    "0201/astro-ph0201038_arXiv.txt": {
        "abstract": "We describe the results of a survey of the UV absorption properties of the Boroson \\& Green sample of AGN, which extends from the Seyfert ($M_V\\simeq -21$) to the luminous quasar ($M_V\\simeq -27$) level. The survey is based mostly on \\HST\\ archival data available for $\\gtsim 1/2$ of the 87 sample objects. Our main result is that soft X-ray weak quasars (SXWQs, 10 AGN with $\\alpha_{ox}\\le -2$) show the strongest UV absorption at a given luminosity, and their maximum outflow velocity, $v_{\\rm max}$, is strongly correlated with $M_V$ ($r_S=-0.95$). This suggests that $v_{\\rm max}$ is largely set by the luminosity, as expected for radiation-pressure driven outflows. Luminous SXWQs have preferentially low [\\oiii] luminosity, which suggests they are physically distinct from unabsorbed AGN, while non-SXWQs with UV absorption are consistent with being drawn from the unabsorbed AGN population. We also find an indication that $v_{\\rm max}/v_{\\rm BLR}$ increases with $L/L_{\\rm Edd}$, as expected for radiation-pressure driven outflows. This relation and the  $v_{\\rm max}$ vs. $M_V$ relation may indicate that the radiation-pressure force multiplier increases with luminosity, and that the wind launching radius in non-SXWQs is $\\sim 10$ times larger than in SXWQs. ",
        "introduction": "Fast outflows are common in AGN. The outflow properties are very different in high and low-luminosity AGN. Whereas luminous high-$z$ quasars display outflows reaching a few $10^4$~km~s$^{-1}$ in $\\sim 10$\\% of the objects, Seyfert galaxies display typical outflow velocities up to only $\\sim 10^3$~km~s$^{-1}$, but in $\\sim 50$\\% of the objects. {\\em Why are the outflow properties so different at low and high-luminosity?  How do the outflow properties vary with luminosity? Are AGN with outflows normal AGN seen at a preferred angle, or are they physically distinct objects? How is the UV absorption related to X-ray absorption?} The answers bear important clues to the origin and acceleration mechanism of AGN outflows.  The largest and most systematic study of broad absorption line quasars (BALQs) was carried out by Weymann et al. (1991), using ground-based spectroscopy of $z>1.5$ quasars. More recent well-defined samples of BALQs are emerging from other large ground-based surveys (Becker et al. 2000; Menou et al. 2001; Hall et al. 2001). However, these comprehensive studies are limited to luminous quasars. Studies of absorption in low-luminosity AGN are best done in low-$z$ objects and thus require \\HST\\ (see Crenshaw et al. 1999 for the most comprehensive study). There are many detailed studies of various objects ranging in luminosity from Seyferts to quasars, but these are focused on understanding the absorber properties of each object and do not provide a clear comprehensive picture.  In this study we make a first step toward a uniform and systematic survey of the UV absorption properties of AGN from the Seyfert to the luminous quasar level. We primarily use archival \\HST\\ observations of the Boroson \\& Green (1992; hereafter BG92) sample, which includes the 87 $z<0.5$ Palomar-Green (PG, Schmidt \\& Green 1983) AGN, and extends from Seyferts at $M_V\\simeq -21$ to luminous quasars at $M_V\\simeq -27$. The high-quality optical data set available from BG92 allows us to explore relations between quasar absorption and emission properties known as ``eigenvector 1'' (EV1; BG92; Boroson 2002),\\footnote {EV1 is defined in BG92, and represents a set of correlated optical emission properties, including in particular the strength of the [\\oiii] and Fe~II emission, and the \\Hbeta\\ line width.} which are also strongly correlated with X-ray and UV emission properties (e.g., Laor et al. 1997, hereafter L97; Brandt \\& Boller 1999; Wills et al. 1999; Laor 2000). These correlations are particularly interesting since they may be driven by fundamental parameters, specifically the black hole mass and the accretion rate.  The UV absorption properties of some AGN from the BG92 sample have been studied by various authors (see the references in Appendix A and in Brandt, Laor, \\& Wills 2000; hereafter BLW, \\S\\S 5, 6), and a complete census of the \\civ\\ absorption equivalent width, \\CEW, of all BG92 AGN with sufficient quality UV spectra is provided by BLW. This study also provides a complete systematic census of the effective optical-to-X-ray spectral slope, $\\alpha_{ox}$, based on \\ROSAT\\ data.\\footnote{Note that more recent \\HST\\ and \\ASCA\\ observations of some of the objects indicate that in a few cases $\\alpha_{ox}$ may vary significantly (e.g., Gallagher et al. 2001).} The $\\alpha_{ox}$ was used by BLW to define a complete sample of Soft X-ray Weak quasars (SXWQs), which includes all 10 AGN from BG92 with $\\alpha_{ox}\\le -2$. The $\\alpha_{ox}\\le -2$ cutoff is based on the distribution of $\\alpha_{ox}$ in the BG92 sample, which suggests an apparently distinct group of SXWQs (see BLW). Since $\\langle \\alpha_{ox}\\rangle\\simeq -1.48$ in non-SXWQs (Laor et al. 1997, and \\S 3.2 here), the 2~keV luminosity of SXWQs is suppressed by a factor of $\\ge 25$ on average, compared to non-SXWQs with the same optical luminosity. BLW discovered a very strong relation between $\\alpha_{ox}$ and \\CEW, which suggests that soft X-ray weakness is due to absorption, as was directly confirmed for some objects by hard X-ray spectroscopy (e.g., Gallagher et al. 1999, 2001; Green et al. 2001).   The \\CEW\\ of the 10 SXWQs covers a wide range ($1-100$~\\AA), and although nearly all BALQs are SXWQs (e.g., Kopko, Turnshek, \\& Espey 1994; Green et al. 1995; Green \\& Mathur 1996), the converse is not true as only some of these SXWQs can be defined as BALQs (see below). {\\em What determines if a SXWQ is a BALQ?} The purpose of this paper is to answer this question, together with some of the questions posed in the opening paragraph. The paper is organized as follows. In \\S 2 we briefly describe the method of analysis, and in \\S 3 we provide the absorption statistics and discuss their dependence on \\MO3, $M_V$, radio-loudness, emission-line strengths, and fundamental parameters. Some further open questions are discussed in \\S 4, and the main conclusions are summarized in \\S 5. Appenix A provides notes on individual objects, supplementing the notes in BLW, and Appendix B lists the 28 objects without intrinsic UV absorption. ",
        "conclusions": "We describe the results of a study of the \\civ\\ absorption properties of AGN extending from the Seyfert level ($M_V\\simeq -21$) to the luminous quasar level ($M_V\\simeq -27$). The study is based on spectra of an incomplete subset of 56 AGN from the 87 BG92 AGN, mostly extracted from the \\HST\\ archives. The main results are the following:  \\begin{enumerate} \\item About 10\\% of the objects show strong (EW$>$10~\\AA) \\civ\\ absorption, $\\sim 20$\\% show intermediate-strength (1~\\AA--10~\\AA) absorption, $\\sim 20$\\% show weak (0.1~\\AA--1~\\AA) absorption, and $\\sim 50$\\% show no ($<$0.1~\\AA) absorption.  \\item SXWQs are $\\sim10$ times more common at \\MO3$>-27$ compared to \\MO3$<-27$, but the frequency of UV absorption in objects without strong X-ray absorption is independent of \\MO3. This indicates that UV absorption without strong soft X-ray absorption may be purely an inclination effect, but apparently higher column UV + soft X-ray absorption (yet without much optical extinction), as also seen in BALQs, occurs in a physical component present only in certain types of AGN.  \\item SXWQs have a higher \\CEW\\ and $v_{\\rm max}$ at a given $M_V$ than UV-only absorbed AGN, and both parameters in SXWQs  are strongly correlated with $M_V$ ($r_S=-0.93$ and $-0.95$). Thus, luminous SXWQs are BALQs, and lower luminosity SXWQs have ``mini-BALs'' or ``associated absorbers''. The observed dependence $v_{\\rm max}\\propto L^{0.62\\pm 0.08}$ is steeper than expected for constant-$\\Gamma$, radiation-pressure driven outflows launched at $R\\sim R_{\\rm BLR}$. This may indicate that $\\Gamma$ increases with luminosity.  \\item There is an indication that the relative outflow velocity, $v_{\\rm max}/v_{\\rm BLR}$, increases with $L/L_{\\rm Edd}$ for all AGN, and that the wind in non-SXWQs is launched at $\\sim 10$ times the launching radius in SXWQs.  \\end{enumerate}  The results presented here are based on UV spectra of 56 of the BG92 AGN, a few of which have only low-quality spectra. To establish the strength and statistical significance of these results it is important to obtain high-quality UV coverage of the complete BG92 sample of 87 AGN. This survey, together with the BLW study, will establish complete UV + soft X-ray coverage of the sample. This can then be followed by detailed UV and X-ray spectroscopy of all the ``interesting'' AGN, which may lead to solutions of some of the above questions. In particular, the class of ``intermediate'' UV and X-ray absorption AGN (seven identified here, Fig.~2) may be especially suitable for studying the relation of the UV and X-ray absorbers.  This first complete UV survey of low-$z$ AGN, together with ground-based surveys of high-redshift AGN samples, will also allow one to explore if there is significant evolution in AGN absorption properties with $z$. A complete survey is a resource-intensive approach but, unlike the study of individual AGN, will allow one to draw conclusions about the general optically selected AGN population. In addition, it is important to define a significantly larger complete sample of SXWQs (e.g., Risaliti et al. 2001) to test the strength and significance of the relations described above.  We thank Todd Boroson for providing us with accurate redshifts for all objects in the BG92 sample, and Bev Wills and Sarah Gallagher for their helpful comments. We also thank the referee for a very detailed and helpful report. This research was supported by grant \\# 209/00-11.1 of the Israel Science Foundation to A.L. and NASA LTSA grant NAG5-8107 to W.N.B.  \\appendix \\begin{center}A. NOTES ON INDIVIDUAL OBJECTS\\end{center}  Relevant notes on some of the objects appear in BLW (\\S\\S 5,6). Here we provide additional notes, mostly on objects with weak \\civ\\ absorption.  0003+158-- BLW note \\CEW=0.8~\\AA\\ for this object, which refers to the $z=0.366$ metal absorber noted by Jannuzi et al. (1998, hereafter J98). Since the associated \\Lya\\ line is unresolved (FWHM$<250$~\\kms), the velocity shift of the system is very large ($v=-17,500$~\\kms), and the higher order Lyman series lines indicate a (partially) optically thin absorber, this is likely to be an intervening system and we adopt here \\CEW=0~\\AA. However, we cannot rule out an intrinsic narrow high-velocity system, as suggested in some high-redshift quasars (e.g., Jannuzi et al. 1996; Hamann, Barlow, \\& Junkkarinen 1997; Richards et al. 1999), though the absorption in the more secure such cases is significantly broader than 250~\\kms.  0007+106-- \\HST\\ did not observe \\civ. There may be very weak \\nv\\ absorption (EW$\\sim 0.2$~\\AA), as also suspected by Crenshaw et al. 1999 (\\S B1 there). The \\Lya\\ EW used here is likely an overestimate of the \\CEW. No significant X-ray absorption is seen with \\ROSAT\\ (Wang et al. 1996).  0050+124-- Weak absorption in \\Lya, \\nv, and \\civ\\ was noted by L97 and Crenshaw et al. (1999). A closer inspection reveals also the \\siIV$\\lambda\\lambda$1393.8,1402.8 doublet in absorption at the same velocity shift of $-1850$~\\kms. Small excess absorption may be present in soft X-rays (Boller, Brandt \\& Fink 1996).  0844+349-- \\HST\\ did not observe \\civ. Wang et al. (2000) noted the $v\\sim 0$~\\kms\\ absorption system in \\Lya. The \\nv\\ absorption is weak (EW$\\sim 0.1$~\\AA), but both doublet components appear to be present. Wang et al. suggest in addition that the absorption at observed-frame 1260.4~\\AA\\ (rest-frame $-7770$~\\kms) is too strong and too broad to be just due to Galactic \\siII$\\lambda 1260.4$ absorption and conclude there is intrinsic \\Lya\\ absorption in this feature as well. However, the Galactic \\cii$\\lambda 1335$ absorption has a very similar profile (see Fig.~3), which suggests the 1260.4~\\AA\\ feature is pure Galactic absorption. There is also evidence for a small excess ``cold'' X-ray absorbing column density in an \\ASCA\\ observation (George et al. 2000, hereafter G00). Note that this object shows large amplitude changes in $\\alpha_{ox}$ (Gallagher et al. 2001, Fig.~9 there).  0923+201-- \\HST\\ did not observe \\civ, and only part of the blue wing of \\nv\\ was observed. However, the ``valley'' between \\nv\\ and \\Lya\\ is much deeper than normally seen in AGN spectra, where the blue wing of \\nv\\ generally blends into \\Lya\\ with only a slight drop in flux density (see all the other profiles in Figs. 1-3). The bottom of this dip at $v\\sim -2500$~\\kms\\ matches quite well the \\Lya\\ absorption system at $-2500$~\\kms. The absorption in both \\Lya\\ and \\nv\\ appears to include a broad, shallow component ($-4000\\ltsim v\\ltsim -2000$~\\kms) on top of a narrow unresolved core.  0947+396-- BLW suggest weak \\civ\\ absorption (EW$\\sim 0.2$~\\AA) based on the line-peak asymmetry. Since there is no clear absorption dip in the line profile we adopt a more conservative estimate of EW=0~\\AA. \\ROSAT\\ and \\SAX\\ spectra do not find significant absorption (L97; Mineo et al. 2000, hereafter M00).   0953+414-- J98 did not note intrinsic \\civ\\ absorption in this object, but Ganguly et al. (2001) did. The \\civ\\ absorption is only marginally significant, though both doublet components seem to be present. The velocity shift of $-1380$~\\kms\\ found by Ganguly et al. is based on $z=0.239$. The revised value of $z=0.23405$ based on the peak of [\\oiii]$\\lambda 5007$ (see Table~1) gives a velocity shift of only 100~\\kms. \\ROSAT\\ and \\ASCA\\ spectra (L97; G00) suggest there may be a small intrinsic X-ray absorbing column density.  1011$-$040-- BLW give \\CEW$<$1.2~\\AA\\ based on the available \\IUE\\ spectrum. However, the marginally significant absorption centered at $v= -200$~\\kms\\ in \\civ\\ appears to be present in \\nv\\ and \\Lya\\ as well. We therefore adopt \\CEW$=1$~\\AA\\ for this object, though a higher quality spectrum is obviously required to verify this result.  1049$-$005--  The \\civ\\ absorption was noted by J98 and by Ganguly et al. (2001). The absorption is very weak, but both doublet components are clearly present.  1100+772-- The \\civ\\ absorption was noted by J98 and by Ganguly et al. (2001). The absorption is rather weak but clearly present in \\civ, \\Lya, and \\ovi. \\ROSAT\\ and \\ASCA\\ spectra do not reveal any intrinsic absorption (Wang, Brinkmann, \\& Bergeron 1996; Sambruna et al. 1999).  1115+407-- The \\civ\\ absorption is weak, but both doublet components are clearly present. The \\nv\\ doublet is also clearly present. \\ROSAT\\ and {\\it BeppoSAX} do not reveal significant intrinsic absorption (L97; M00).  1116+215--  BLW suggest very weak \\civ\\ absorption (EW$\\sim 0.1$~\\AA) based on a feature at $v= -9500$~\\kms. This feature is below the detection threshold of J98, and the second doublet component is not seen. If real, this system is likely to be associated with the intervening \\Lya\\ system at $v= -9500$~\\kms (related to an intervening galaxy; Tripp, Lu \\& Savage 1998), and we therefore adopt here \\CEW=0~\\AA. A very high S/N \\ROSAT\\ spectrum puts a strong constraint on any cold absorber (L97), and an \\ASCA\\ spectrum suggests a weak warm absorber (G00).  1202+281-- BLW give \\CEW=0.4~\\AA. This is based on the \\civ\\ line-peak asymmetry. However, inspection of the \\Hbeta\\ profile in BG92 reveals the same asymmetry. Since \\Hbeta\\ absorption is extremely rare, we conclude that the emission-line asymmetry is most likely not due to absorption, and adopt here \\CEW=0~\\AA. A rather high S/N \\ROSAT\\ spectrum does not reveal any evidence for absorption (L97)  1211+143-- BLW give \\CEW=0.5~\\AA. This is based on marginally significant features at $v\\sim -8000$~\\kms. The two apparent narrow \\civ\\ components are somewhat too close (400~\\kms\\ apart, instead of 500~\\kms) and are displaced by $\\sim 200$~\\kms\\ from the nearest \\Lya\\ absorber. We conclude that the identification of this system is not secure and adopt \\CEW=0~\\AA. High S/N \\ROSAT\\ and \\ASCA\\ spectra do not reveal any intrinsic absorption (Fiore et al. 1994; G00).  1322+659-- BLW give \\CEW=0.2~\\AA, based on a narrow absorption feature at $v=120$~\\kms. A similarly shaped absorption feature occurs in \\Lya, but it is centered at $v=-20$~\\kms. The apparent \\civ\\ absorption is centered at 1808.2~\\AA, and it may be due to Galactic \\siII$~\\lambda 1808.0$ since the other resonance \\siII\\ lines at 1190.4~\\AA, 1193.3~\\AA, 1260.4~\\AA, 1304.4~\\AA, and 1526.7~\\AA\\ are unusually strong. We therefore conservatively adopt \\CEW=0~\\AA. A \\ROSAT\\ spectrum suggests a low-column density intrinsic neutral absorber, and an \\ASCA\\ spectrum suggests a weak warm absorber (L97; G00).  1352+183-- BLW give \\CEW=0.5~\\AA\\ based on a narrow absorption feature at $v=-9900$~\\kms. However, since the second \\civ\\ doublet component does not appear to be present, and the nearest possible \\Lya\\ absorption is at $v=-10200$~\\kms, we conclude that this feature is most likely not due to \\civ\\ and adopt \\CEW=0~\\AA. No absorption is seen with \\ROSAT\\ and \\SAX\\ (L97; M00).  1402+261-- Although \\civ\\ absorption is clearly seen at $v=-4600$~\\kms, two earlier \\HST\\ observations in 1993 (PI Tytler) and in 1994 (Turnshek et al. 1997) do not show significant absorption in \\civ, providing strong evidence that the absorption is variable. No absorption is seen with \\ROSAT\\ and \\SAX\\ (L97; M00).  1427+480-- BLW give \\CEW=0.03~\\AA\\ based on two very weak features at $v=-100$~\\kms\\ and $v=400$~\\kms, and a weak \\Lya\\ absorption feature at $v=0$~\\kms. Since the amplitudes of these features are comparable to the noise we adopt here \\CEW=0~\\AA. No absorption is seen with \\ROSAT\\ (L97).  1512+370-- BLW give \\CEW=0.2~\\AA\\ based on two apparently significant features at $v=250$~\\kms\\ and $v=800$~\\kms. However, since there is no clear corresponding \\Lya\\ feature we adopt \\CEW=0~\\AA\\ here. J98 and Ganguly et al. (2001) also do not identify these features as \\civ\\ absorption. No absorption is seen with \\ROSAT\\ and \\SAX\\ (L97; M00).  1543+489-- BLW give \\CEW=0.4~\\AA\\ based on an apparently very broad and very shallow absorption dip at the \\civ\\ line peak ($v=-1000$~\\kms\\ to $-3000$~\\kms). However, since we cannot rule out a peculiar line profile in this rather extreme EV1 object (BG92), we adopt \\CEW=0~\\AA\\ here.  2214+139--  BLW give \\CEW$<1.2$~\\AA\\ based on the available \\IUE\\ spectrum. However, the peak of \\nv\\ appears to be redshifted by $\\sim 1000$~\\kms, which is highly unusual and may imply absorption of the blue wing of \\nv. Some absorption may be present in \\civ\\ at $v=-1000$~\\kms, though \\Lya\\ shows no evidence for absorption. We tentatively estimate a marginally significant \\CEW=1.1~\\AA\\ and use a different symbol for this object throughout the paper.  2251+113-- BLW give \\CEW=0.8~\\AA\\ based on a very deep resolved \\civ\\ doublet at $v=0$~\\kms. Here we identify additional absorption extending up to $v=-5000$~\\kms, which increases \\CEW\\ to 3.5~\\AA\\ (see also the comment in \\S 6.4 of Ganguly et al. 2001). The velocity shift of $610$~\\kms\\ found by Ganguly et al. is based on $z=0.323$. The revised value of $z=0.32553$ (see Table~1) gives a velocity shift of 0~\\kms.  \\appendix \\begin{center}B. OBJECTS WITH \\CEW$<0.1$~\\AA \\end{center}  For the sake of completeness we list the 28 objects where we find intrinsic \\CEW$<0.1$~\\AA. This limit is based on \\civ, \\Lya, and \\nv. In objects where these three lines are not all available we list in parentheses the line which was available and used to constrain the UV absorption. The ``no UV absorption'' sample includes the following objects: PG~0003+158, PG~0003+199, PG~0026+129, PG~0052+251 (\\civ), PG~0947+396, PG~1103$-$006, PG~1116+214, PG~1121+422 (\\civ), PG~1202+281, PG~1211+143, PG~1216+069, PG~1226+023, PG~1259+593, PG~1302$-$102, PG~1307+085 (\\civ), PG~1322+659, PG~1352+183, PG~1415+451, PG~1416$-$129, PG~1427+480, PG~1440+356, PG~1444+407, PG~1512+370, PG~1534+580 (\\Lya), PG~1543+489, PG~1545+210, PG~1612+261 (\\civ), PG~1626+554."
    },
    "0201/astro-ph0201381_arXiv.txt": {
        "abstract": "A short history is given of the development of the correction for observation selection bias inherent in the calibration of absolute magnitudes using trigonometric parallaxes. The developments have been due to Eddington, Jeffreys, Trumpler and Weaver, Wallerstein, Ljunggren and Oja, West, Lutz and Kelker after whom the bias is named, Turon Lacarrieu and Cr\\'{e}z\\'{e}, Hanson, Smith, and many others. As a tutorial to gain an intuitive understanding of several complicated trigonometric bias problems, we study a toy bias model of a parallax catalog which incorporates assumed parallax measuring errors of various severities. The two effects of bias errors on the derived absolute magnitudes are (1) the Lutz-Kelker correction itself that depends on the relative parallax error $\\delta \\pi / \\pi$ and the spatial distribution, and (2) a Malmquist-like `incompleteness' correction of opposite sign due to various apparent magnitude cut-offs as they are progressively imposed on the catalog. We calculate the bias properties using simulations involving $3 \\times 10^{6}$ stars of fixed absolute magnitude using $M_{v} = +0.6$  to imitate RR Lyrae variables in the mean. These stars are spread over a spherical volume bounded by a radius 50,000 parsecs with different spatial density distributions. The bias is demonstrated by first using a fixed rms parallax uncertainty per star of $ 50 \\muas$ , and then using a variable rms accuracy that ranges from $ 50 \\muas $ at apparent magnitude $V = 9$ to $500 \\muas$ at $V = 15$ according to the specifications for the {\\it FAME} astrometric satellite to be launched in 2004. The effects of imposing magnitude limits and limits on the `observer's' error, $\\delta \\pi / \\pi$, are displayed. We contrast the method of calculating mean absolute magnitude directly from the parallaxes where bias corrections are mandatory, with an inverse method using maximum likelihood which is free of the Lutz-Kelker bias, although a Malmquist bias is present. Simulations show the power of the inverse method. Nevertheless, we recommend reduction of the data using both methods. Each must give the same answer if each is freed from systematic error. Although the maximum likelihood method will, in theory, eliminate many of the bias problems of the direct method, nevertheless the bias corrections required by the direct method can be determined {\\it empirically} via Spaenhauer diagrams immediately from the data, as discussed in the earlier papers of this series.  Any correlation of the absolute (trigonometric) magnitudes with the (trigonometric) distances {\\it is the bias}. We discuss the level of accuracy that can be expected in a calibration of RR Lyrae absolute magnitudes from the {\\it FAME} data over the metallicity range of $ {\\rm [Fe/H]} $ from $0$ to $-2$, given the known frequency of the local RR Lyraes closer than 1.5 kpc.  Of course, use will also be made of the entire {\\it FAME} database for the RR Lyrae stars over the complete range of distances that can be used to empirically determine the random and systematic errors from the {\\it FAME} parallax catalog, using correlations of derived absolute magnitude with distance and position in the sky. These bias corrections are expected to be much more complicated than only a function of apparent magnitude because of various restrictions due to orbital constraints on the space-craft. ",
        "introduction": "The approved NASA mission, {\\it FAME}, is a science program based on an astrometric satellite that will obtain all-sky trigonometric parallaxes and proper motions for stars brighter than magnitude 15. The accuracy of the parallaxes has been specified to be $ 50 \\muas$  in the best range of the space-craft's configuration for stars brighter than $V = 10$ mag, degrading to no worse than $500 \\muas$ at its detection limit at $V = 15$. These accuracies are between 2 and 20 times more accurate than achieved by Hipparcos. The typical rms accuracy for trigonometric parallaxes of Hipparcos data is $1000 \\muas$, spectacular at the time but not accurate enough by at least a factor of 10 to reach the domain of the RR Lyrae stars. The promise of the data from {\\it FAME}, if its mission goals are achieved, is that the domain needed for the RR Lyrae absolute magnitude calibration can be achieved directly via trigonometric parallaxes.  The purpose of the present paper is to assess what must be done with the RR Lyrae parallax data from {\\it FAME} in order to correct for observational selection bias in determining a correct calibration of absolute magnitude for these stars as a function of their metallicity.  It is widely recognized that (1) the solution to many problems in Galactic structure, (2) an account of the episodes and time-sequences in the formation of the Galaxy, and (3) one approach to the extra-galactic distance scale, rest directly on the calibration of $M_{V}(RR) = f({\\rm [Fe/H]})$. In particular, the steepness of the dependence of $M_{V}(RR)$ on ${\\rm [Fe/H]}$ determines whether there is an appreciable time interval over which the Galactic globular clusters of different metallicities have formed, or whether the entire Galactic globular cluster system formed nearly simultaneously with the rapid collapse of the nascent Galaxy with its early separation of the disk and the halo (Baade 1957; Eggen et al. 1962; Sandage 1986, 1990a).  It has been shown elsewhere (Sandage and Cacciari 1990) that if the slope of the relation between absolute magnitude and metallicity for RR Lyrae stars is as large as $dM_{V}(RR)/d([Fe/H]) = 0.32$, then there is no dependence of the ages of the Galactic globular clusters on $[Fe/H]$. This conclusion was based on the stellar models available in 1990, before the $[O/Fe]$ enhancement was known. However, the same dependence of the globular cluster formation history on the value of the $dM_{V}(RR)/d([Fe/H])$ slope has more recently been confirmed by Chaboyer, Demarque, and Sarajedini (1996) using Oxygen enhanced models.  The dependence of the age spread on the slope was not discussed by Chaboyer {\\it et al.}  They only treated the $dM_{V}(RR)/d([Fe/H]) = 0.20$ case, but their large age spread for this low-slope case is nearly identical with that of Sandage and Cacciari (1990) for the same small slope. If they had used the steeper metallicity dependence that is required by the observed Oosterhoff period effect (Sandage 1993a,b), their conclusion concerning a large age spread depending on $[Fe/H]$ would have been reversed. No age spread is predicted if the preferred slope of 0.30 is used. In fact, the Oxygen enhanced models of Bergbush and Vandenberg (1992) show that if the slope is as shallow as $dM_{V}((RR))/([Fe/H]) = 0.26$, then there is no age spread among the globular clusters of different metallicities. Clearly, one crucial importance of the {\\it FAME} mission will be its ability to determine a definitive calibration of $ M_{V}(RR) = f([Fe/H]) $ relation for RR Lyrae stars. The data impact directly on the formation history of the Galaxy.  However, impressive as  the specified accuracy of $50 \\muas$ is for the parallax accuracy for stars at $V = 10$ mag, that accuracy is near the margin of what is needed to make a definitive calibration of $M_{V}(RR)$ as a function of metallicity. The uncertainties center on the inevitable observational selection bias due to the distribution of parallax errors that will exist in the highly non-linear distribution of the observed parallaxes. This effect is currently named the Lutz-Kelker bias. The problem has a long history, part of which we review in the next section. ",
        "conclusions": ""
    },
    "0201/astro-ph0201348_arXiv.txt": {
        "abstract": "We present submillimeter polarimetry at 850 \\micron\\ toward the filamentary star-forming region associated with the reflection nebulosity NGC 2068 in Orion B.  These data were obtained using the James Clerk Maxwell Telescope's SCUBA polarimeter.  The polarization pattern observed is not consistent with that expected for a field geometry defined by a single mean field direction.  There are three distinct distributions of polarization angle, which could represent regions of differing inclination and/or field geometry within the filamentary gas.  In general, the polarization pattern does not correlate with the underlying total dust emission.  The presence of varying inclinations against the plane of the sky is consistent with the comparison of the 850 \\micron\\ continuum emission to the optical emission from the Palomar Optical Sky Survey, which shows that the western dust emission lies in the foreground of the optical nebula while the eastern dust emission originates in the background. Percentage polarizations are high, particularly toward the north-east region of the cloud.  The mean polarization percentage in the region is 5.0\\% with a standard deviation of 3.1\\%.  Depolarization toward high intensities is identified in all parts of the filament. ",
        "introduction": "The study of polarized emission from molecular clouds is of great interest since, at long wavelengths ($\\lambda > 25$ \\micron), it effectively traces the orientation on magnetic fields local to star-forming regions \\citep{hil88}.  Magnetic fields have been shown to be energetically comparable to gravity and kinetic motions within molecular clouds \\citep{mg88,cru99,basu00} and are theorized to provide vital support to clouds, preventing global collapse \\citep[and references therein]{mou87,mck93}.  Such support is necessary to explain the low star-forming efficiencies observed in molecular clouds, including Orion B, where the efficiency is $\\sim 1$\\% \\citep{car00}.  The Orion B cloud, at a distance of $\\sim 415$ pc \\citep{ant82}, is the closest giant molecular cloud.  Within it, star formation is concentrated into five distinct regions: NGC 2071, NGC 2068, LBS 23 (HH 24), NGC 2024 and NGC 2023, as determined from unbiased surveys for young stellar objects \\citep{lad91a} and dense CS gas \\citep{lad91b}.  Large scale 850 \\micron\\ dust emission from these regions has been mapped by \\citet{mit01} and \\citet{mot01}.  Maps of the polarized emission from NGC 2071, LBS 23, and NGC 2024 have already appeared in the literature (see \\citet[hereafter Paper III]{mat01}).  In order to compare field geometries across Orion B, we have now mapped the polarized emission from NGC 2068.  The submillimeter emission from NGC 2068 lies south of the reflection nebula (see Figure \\ref{p4:pass+scuba}) which is seen optically and contains an infrared cluster \\citep{lad91a}.  The overall structure of the dust emission is that of a ``clumpy filament'' in which 18 distinct compact sources are identified.  Most of the submillimeter sources fall outside the boundary of the cluster identified in CS by \\citet{lad91b}.  Star formation is ongoing in this region, as evidenced by detections of bipolar outflows by \\citet{mit01} around OriBN 51, from which evidence of outflow previously existed \\citep{edw84,gib00}, and also around OriBN 35 and OriBN 36 peaks (see their Fig. 1c or Figure \\ref{p4:map} below), although the sources of these outflows are ambiguous due to the close proximity of OriBN 33 and OriBN 37 to the positions of high velocity gas.  OriBN 39 has a 2MASS infrared source associated with it, and thus should be a source of outflow.  Evidence for redshifted gas near OriBN 47 also suggests some outflow from this source.  There are no prior observations of polarized emission from the NGC 2068 dust emitting region.  However, \\citet{man84} measured linear polarization from scattering against the reflection nebula north of the molecular condensations.  Based on the centrosymmetric pattern observed, \\citet{man84} ruled out the presence of aligned grains within the reflection nebula.  They infer the presence of a foreground assembly of grains illuminated from behind solely by the star HD 38563N.  The position of this star is $\\sim 5$\\arcmin\\ north-east of the dust emission on which we report in this work, concident with the near-IR cluster observed by \\citet{lad91a}.  Polarization data probe only two directions of the magnetic field geometry -- those on the plane of the sky.  Additionally, they provide no information about the strength of the magnetic field.  Where field geometries are simple and the direction of the magnetic field does not vary through the cloud depth, the polarized emission detected is perpendicular to the mean magnetic field and the latter can be inferred simply by rotating the polarization vectors by $90^\\circ$. If the field has a more complex, non-uniform geometry, then interpretation becomes more difficult.  In such cases, it is best to compare directly the polarization maps with polarization patterns predicted from a physical model of a magnetized cloud.  It is important to recognize that the polarizations measured are vector sums along a particular line of sight through the cloud observed, weighted by column density.  \\citet{fp00a} present a model for a filamentary cloud in which a helical magnetic field threads the filament and plays an important role in determining the radial density structure.  This model predicts an $r^{-2}$ density profile, which has been observed in several clouds, including the \\intfil\\ \\citep{joh99} and several clouds in Cygnus \\citep{lal99,all99}. The helical field geometry also predicts depolarization toward the axis of a filament due to cancellation effects on either side of the axis.  \\citet{fp00c} present predicted polarization patterns for cases in which the field is either poloidally or toroidally dominated as well as for various filament inclinations.  Qualitative extensions to these models have been shown to reproduce observed polarization patterns in the filamentary clouds OMC-3 \\citep[hereafter Paper II]{mwf01} and NGC 2024 (see Paper III).  The filamentary structure of NGC 2068 is therefore of particular interest as a further test of axially symmetric magnetic field geometries.  This paper is the fourth in a series which seeks to compare the polarization patterns (and inferred magnetic field geometries) in different star-forming regions.  The observations and data reduction techniques are described in $\\S$\\ref{p4:obs}.  The polarization data are presented and analyzed in $\\S$\\ref{p4:poldata}.  The implications of these data for the local magnetic field geometry and that of Orion B as a whole are discussed in $\\S$ \\ref{p4:disc}, and $\\S$ \\ref{p4:summ} summarizes our results. ",
        "conclusions": "\\label{p4:disc}  \\subsection{Polarization Percentage in Orion B}  Figure \\ref{p4:Pcompare} shows the distributions of polarization percentage in the regions of Orion B observed with the SCUBA polarimeter.  These distributions have been normalized to the total number of vectors measured in each region, and the counts expressed as a fraction of the total.  The error bars represent the $\\sqrt{N}$ statistical uncertainty in the number of counts, also normalized to the total in each map.  The distribution for NGC 2068 is comparable to those measured in the other star-forming regions of Orion B North -- NGC 2071 and LBS23N -- and with measurements toward the OMC-3 filament in Orion A (Paper II). The NGC 2024 region is dominated by low percentage polarizations, with the majority between 1-2\\%.  This result could imply NGC 2024 has weaker magnetic fields, poorer grain rotation or alignment, a different grain composition, or some combination of these factors. The fact that NGC 2071, NGC 2068 and LBS 23N show distributions which are reasonably flat from 1-6\\% could indicate similar grain properties, field strengths, and degrees of grain aligment, but this cannot be proven with these data.  The densities toward the NGC 2024 cores have been estimated to be abnormally high ($10^8$ cm$^{-3}$ \\citet{mez88}), but \\citet{sch91} suggest the values are more typical of cores ($10^6$ cm$^{-3}$).  Polarization data at other wavelengths (such as 350 \\micron) could help constrain the dust properties according to models (see \\citet{hil00} and references therein), and observations of dense molecular tracers like OH for Zeeman splitting could provide more detailed information about the field geometries in these four regions.  Toward all four regions of Orion B observed in polarized emission, the depolarization effect is detected (see Fig.\\ 5 of Paper III and Figure \\ref{p4:depol}).  Paper III shows that NGC 2024 exhibits a similar depolarization signature to the northern core NGC 2071.  However, LBS 23N has a significantly steeper slope of $-0.95$, which is more consistent with that of Region 3 in our data set.  This trend has been observed in many other regions as well, including massive cores such as OMC-1 \\citep{sch97} and protostellar and starless cores \\citep{gir99,war00}.  \\subsection{Evidence for Varying Inclinations in NGC 2068}  The filamentary dust emission of NGC 2068 likely arises from dust at different depths in the Orion B cloud.  The comparison between the Palomar Observatory Sky Survey optical data and the dust emission has been made by \\citet{mit01}. We produce a similar image in Figure \\ref{p4:pass+scuba}, which shows that while the western dust emission lines up well with the optically dark dust lane, the eastern dust emission has no corresponding dark lane.  (North of the \\hii\\ region lies another dust condensation not yet observed with the polarimeter.) Hence, it is likely that the dust emission may lie on the outer edge of the reflection nebula with the western material in the foreground and the eastern material behind.  Thus, it is clear that this filament does not lie in the plane of the sky, and that the inclination on the sky is likely variable.  An obvious question is whether or not all the dust emission arises from material which is spatially related.  It is noteworthy that the core OriBN 34 does not appear to line up well with the filamentary material of Figure \\ref{p4:map}.  In fact, our comparison of the optical and dust emission shows that the location of OriBN 34 is optically dark.  This could be a coincidence, or OriBN 34 could be closer than the dust emission appearing next to it on the SCUBA image. OriBN 39 also appears dark and could be in the foreground. \\citet{mit01} present $^{13}$CO $J=2-1$ toward NGC 2068 and a partial map of the north-eastern region in C$^{18}$O $J=2-1$ (see their Figs. 6 and 7).  The $^{13}$CO contours are closely correlated to the 850 \\micron\\ dust emission.  The same is true for the C$^{18}$O emission where data exist.  These maps are integrated over velocity ranges from 5 \\kms\\ $< v_{LSR} <$ 15 \\kms\\ and 7 \\kms\\ $< v_{LSR} <$ 13 \\kms, respectively.  Thus, the emission is confined to the Orion B cloud.  The fact that OriBN 34 and 39 show similar polarization position angles as the eastern area but may be in the foreground would argue against these regions being completely spatially distinct (unless the cores are not contributing as much polarized emission as the background diffuse material.  If the material is spatially separated, then the similar orientations of the polarization position angles could suggest that the field, if not defined by a single mean field direction, could be organized on spatial scales at least as large as the separation between them.  The 2.4 \\micron\\ emission from the NGC 2068 reflection nebula spans approximately 6\\arcmin\\ in spatial extent \\citep{sel84}.  This is consistent with the physical size of the nebula in optical emission as shown in Figure \\ref{p4:pass+scuba}, and at a distance of 415 pc corresponds to 0.7 pc.  If the nebula is as deep as it is wide, then the source of dust emission at the north-east could be more than a parsec displaced spatially from the cores at the west (taking 7\\arcmin\\ as an estimate of their separation in projection).  \\subsection{Field Geometry}  \\subsubsection{A Single Mean Field Direction}  The vectors of NGC 2068 do not support a single field orientation in NGC 2068.  If one assumes that the magnetic field geometry throughout a cloud's depth is reasonably well defined by a mean field direction, then the direction of the field can be obtained by rotating emission polarization data by $90^\\circ$.  Doing this in the NGC 2068 polarization map will clearly not produce aligned vectors, any more than the polarization data themselves are aligned.  Therefore, we can rule out a uniform, unidirectional field across all of NGC 2068.  This was also the case in OMC-3 in Orion A (see Paper II); however, in that region, the filamentary axis was easy to define, and a strong correlation between the axis and polarization position angle was observed along 75\\% of the filament.  The vectors of NGC 2068 do not show an obvious alignment of polarization position angle with the filament orientation.  We can compare the polarization position angles across those regions observed thus far in Orion B.  In the three regions of Orion B North, no evidence for similar polarization orientations exists.  Toward NGC 2071, the polarization vectors align generally with the prominent outflow from the IRS 1 source at a position angle of $40^\\circ$ east of north.  Within the LBS 23N string of cores to the south of NGC 2068, the vectors are generally aligned north-south (position angle $0^\\circ$) although the scatter in this faint region is considerable. Neither of these orientations is dominant in NGC 2068.  Thus, there is no support for a mean field direction in NGC 2068 or across the three star-forming regions of Orion B North.  The one region observed in Orion B South is NGC 2024, and the polarization pattern from that region has been modeled as arising either from a helical field geometry or from the expansion of the ionization front due to the associated \\hii\\ region in Paper III. The latter geometry is favored since it is most compatible with the total physics and geometry of NGC 2024.  It is clear that the polarization patterns across these star-forming regions can be strongly correlated with the dust and gas structures of a particular region, and that each region must be modeled separately.  \\subsubsection{More Complex Geometries}  As discussed above, there is evidence that the dust arises from a connected gas structure.  However, the spatial separation between the filament edges could be $> 1$ pc.  The regions of different position angles could indicate regions of different field geometry or inclinations.  We propose two model geometries which could potentially explain the variable position angles observed.  Since the cores are mainly aligned along a filamentary structure, a helical field geometry is appealing.  It can explain the confinement of gas in the filament, the fragmentation to cores and the elongation of the asymmetric concentrations along the axis of the filament, which is certainly the case for OriBN 31, 32, 33, 41 and 49.  \\citet{fp00c} show possible polarization patterns for different helical field conditions (i.e.\\ poloidally-dominated versus toroidally-dominated). These models are developed for straight filaments, and \\citet{fp00c} find that for such filaments, the polarization vectors align either parallel or perpendicular to the projected filament axis regardless of the inclination of the filament.  However, Paper II presents a model for a helical field threading a bent filament.  In this case, the vectors may adopt any orientation relative to the filament due to the asymmetries in the filament.  The relative alignment depends on the filament's inclination and rotation on the plane of the sky.  Thus, as inclination and angle in the plane of the sky vary, a helically-threaded filament should produce different polarization position angles relative to the projected filament orientation.  For instance, if the polarization positions are aligned with the filament, a toroidally-dominated helical field could wrap the filament locally. Conversely, where the vectors appear perpendicular to the filament, a helical field would be poloidally-dominated, or the filament and field could be significantly bent.  At the north-western edge near the cores OriBN 33 and 41, the vectors are neither parallel nor perpendicular to the filament.  This could also indicate that the filament is bent along its length in this area.  Modeling such a complex filamentary structure will be difficult, since more than one filament/field geometry can produce similar polarization patterns.  One additional test for the presence of a helical field geometry is the radial profile of the total, unpolarized emission.  Using the 850 \\micron\\ map of \\citet{mit01}, a radial profile can be built up by taking several slices across the filament.  We have attempted to confine the cuts to the regions between cores along NGC 2068.  Figure \\ref{p4:slice} shows the profiles through several such slices, taken between the OriBN 44 and 48, between 32 and 39 and between 35 and 36. For the second cut, both sides of the profile were used; for the other two, one side of the slice was discarded due to the presence of extended emission from the filament near OriBN 47 (for the cut between OriBN 44 and 48) and extended emission from core OriBN 34 (for the cut between OriBN35 and 36).  The fluxes have been normalized to the maximum through the cuts (respectively 0.07 \\jybeam, 0.32 \\jybeam, and 0.28 \\jybeam); this position is assumed to be the axis of the filament.  The NGC 2068 filament is so narrow and faint, it is difficult to interpret the profiles.  To guide the eye, we have drawn on Figure \\ref{p4:slice} lines corresponding to slopes of -0.5 and -1. Near the axis, the flux falls off with a distribution consistent with a power law index of $-0.5$.  Toward lower fluxes, the index could be closer to $-1$; however, at these low flux levels, interpretation of the index becomes difficult.  The flux profile of $r^{-0.5}$ corresponds to a density profile of $r^{-1.5}$ for an isothermal filament.  The helical field geometry predicts a density profile with index $-2$ for an isothermal equation of state.  For other equations of state (e.g.\\ the logotrope, \\citet{mcl96}), a shallower range of density profiles can be generated.  Unmagnetized, isothermal filaments predict much steeper profiles, with indices of $-4$, which are clearly not consistent even with these poor profiles.  A second possible field geometry is suggested by the position of the molecular filament so near the periphery of the reflection nebula of NGC 2068.  Paper III presents a model of the polarization data toward the NGC 2024 dense ridge of cores in which the field, compressed by the expansion of the NGC 2024 \\hii\\ region, is then moulded around the dense ridge as the ionization front approaches the cores.  This picture accounts for both the polarization pattern at 850 \\micron\\ and the measurements of the line-of-sight field direction and strength measured by \\citet{cru99}.  The expansion of the reflection nebula could produce a similar effect in NGC 2068, although the pattern produced is much more complex.  In NGC 2024, the ridge is entirely located on the far-side of the \\hii\\ region, whereas in NGC 2068 the comparison of optical and 850 \\micron\\ dust emission in Figure \\ref{p4:pass+scuba} indicates that the dust emission arises from both the near and far-side of the nebula.  A recent publication on the formation of quiescent cores through turbulent flows \\citep{pad01} provides a third possible interpretation of the polarization pattern.  This work discusses the potential formation of cores due to the presence of super-sonic turbulent flows within molecular clouds.  In this model, cores form by accretion along filamentary structures, with the brightest cores forming at the loci of intersecting shocks.  The core OriBN 47 lies at the intersection of three filamentary segments.  As predicted by \\citet{pad01}, this core exhibits depolarization, although at the distance of Orion B, it is difficult to achieve a good sampling of polarization vectors across this core.  Thus, we do not see the large changes in position angle predicted near cores (see their Fig.\\ 3) and routinely observed in Bok globules and starless cores in closer regions of star formation (e.g.\\ see \\citet{war00}).  However, the polarization pattern along the larger scale filamentary structure is well sampled in our map.  In the turbulent flows model, the polarization vectors along the filamentary structures are seen to align well with the filaments' axes.  In NGC 2068, only the filament segment east of OriBN 48 exhibits this behavior.  The segments to the north-west and south-west show vectors oriented roughly perpendicular to the filamentary structure, which does not agree well with the turbulent flows picture.  A simulation of turbulent flows along filaments of higher density may agree better with our observations in NGC 2068.  If a threshold in extinction exists beyond which grains are not effectively aligned \\citep{pad01}, then denser turbulent flows could exhibit less correlation between the filamentary axis and the inferred polarization."
    },
    "0201/astro-ph0201313_arXiv.txt": {
        "abstract": "Clusters of galaxies are enriched with positrons from jets of active galactic nuclei (AGNs) or from the interaction of cosmic-rays with the intracluster gas.  We follow the cooling of these positrons and show that their eventual annihilation with cluster electrons yields a narrow annihilation line.  Unlike annihilation in the interstellar medium of galaxies, the line produced in clusters is not smeared by three-photon decay of positronium, because positronium formation is suppressed at the high ($\\ga 1$ keV) temperature of the cluster electrons. We show that if AGN jets are composed of $e^+e^-$ pairs, then the annihilation line from rich clusters within a distance of 100 Mpc might be detectable with future space missions, such as INTEGRAL or EXIST. ",
        "introduction": "Despite decades of intense study, the composition of relativistic radio jets remains enigmatic.  While the existence of synchrotron-emitting electrons is secure in both quasars \\citep{bbr} and microquasars \\citep{mirabel}, there is no conclusive evidence that can determine whether positrons or protons make up the positively charged component.  Naively, one might expect these cases to be easily distinguishable through an observational search for the electron-positron annihilation line.  However, modeling of the annihilation process in active jets has shown this not to be the case: the annihilation spectral feature is not a line but is instead very broad and hence difficult to unambiguously identify (although it may possibly contribute to the observed $\\gamma$-ray spectrum of blazars; see \\citealt{boettcher}).  Even if positron cooling to non-relativistic energies within the jet is efficient, the high bulk Lorentz factor ($\\gamma_{\\rm jet} \\sim 10$) of the jet would likely smear out an annihilation line \\citep{bbr}.  The resulting radiation may be detectable with upcoming space observatories, although its interpretation will depend critically on proper modeling of the region inside the jet \\citep{wang}.  The chief obstacle to producing a recognizable annihilation-line feature is the highly relativistic nature of the associated plasma.  This obstacle may be overcome simply by waiting for the active galactic nucleus (AGN) to become dormant.  Once the central engine disappears, the material in the jet will presumably mix with the ambient medium and cool to the ambient temperature.  In this paper, we calculate the annihilation signal from ``relic'' positrons produced by an AGN embedded in a galaxy cluster, as the positrons thermalize with the electrons in the intracluster medium (ICM). We will show that under reasonable assumptions an annihilation line from nearby clusters would be observable in the near future.  A similar suggestion of thermal annihilation after escape from the accelerating source was made by \\citet{maciolek} in the context of small-scale jets from stellar-mass black holes.  However, in galactic environments with typical temperatures $T \\la 10^6 \\kel$, formation of an annihilation line at $511 \\keV$ is inhibited by the rapid formation of positronium, whose primary annihilation channel yields three photons.  For example, the observations of \\citet{kinzer} show that $93 \\pm 4 \\%$ of pair annihilations in the Galactic center occur through the positronium channel and that $\\ga 70\\%$ of the total annihilation energy is emitted in a broad continuum rather than in the line.  We will argue that galaxy clusters are ideally suited to producing annihilation lines because the characteristic temperature of the ICM is larger than the binding energy of positronium; hence, nearly all annihilations produce photons near $511 \\keV$.  Another important open question involves the dynamical significance of cosmic rays to the ICM.  Faraday rotation studies have revealed that cluster cores have magnetic fields $\\ga 1 \\microgauss$ \\citep{kim,clarke}, while observations of excess emission in the ultraviolet and hard X-ray bands from the ICM have been used to infer the existence of widespread cluster fields at slightly lower levels $\\sim 0.1$--$1 \\microgauss$ (e.g., \\citealt{rephaeli2}).  The pervasiveness of magnetic fields in clusters indicates that collisionless shocks generated by accretion or merger events may be efficient particle accelerators \\citep{bland-eichler,col}. Observations of synchrotron radio halos indicate that the acceleration of electrons by cluster merger shocks is a common occurrence \\citep{kemp-sar}. However, these observations probe only the electron component.  In order to understand the dynamical effects of the cosmic rays on the cluster, we are most interested in the accelerated protons, both because shock acceleration may inject more energy into this component than into the electron component \\citep{fields,but} and because cosmic ray protons do not rapidly lose their energy through radiative cooling.  Recent cosmological simulations suggest that the cosmic ray pressure might be as large as $\\sim 10$--$40 \\%$ of the thermal gas pressure in clusters \\citep{miniati-prot}.  To date, two diagnostics of the cosmic ray proton content have been proposed, both relying on the decay of pions produced in collisions between cosmic ray and thermal protons, namely synchrotron and inverse Compton emission from secondary electrons and positrons produced in charged pion decay \\citep{blasi,dolag,miniati-sec} and $\\gamma$-rays produced in neutral pion decay \\citep{col,miniati-prot}.  The former process suffers from the possibility of contamination by newly accelerated electrons, while the $\\gamma$-ray signal will remain below the detection limits at least until the launch of the GLAST mission\\footnote{See http://www-glast.stanford.edu/mission.html} in 2005, largely because the energy from the decaying pions is distributed over a very wide range of photon energies.  Because of the rapid cooling of positrons produced in the decay of $\\pi^+$ to non-relativistic temperatures, we would expect an annihilation line to be produced in this case as well.  (Here too the formation of positronium is inhibited by the relatively large cluster temperatures.)  We therefore also calculate the annihilation spectrum produced through secondary positron production by cosmic ray protons.  However, we find that this signal is well below the detection thresholds of upcoming instruments. Hence, AGN are the only realistic pollutant of substantial amounts of positrons into galaxy clusters.  We thus argue that the future detection of positron annihilation line radiation from clusters would constitute a robust signature of electron-positron jets.  We begin by describing the factors determining the evolution of the positron population in \\S \\ref{posdfevol}, including cooling, annihilation, and source terms.  We then solve the evolution equations in \\S \\ref{posdfsoln}.  In \\S \\ref{annspec} we calculate the resulting positron annihilation signals.  Finally, we conclude in \\S \\ref{conc} with a discussion of our results and prospects for future observations.  Throughout the paper, we assume a $\\Lambda$CDM cosmology with $\\Omega_0 = 0.3$, $\\Omega_\\Lambda = 0.7$, and $H_0 = 70 \\hunits$. ",
        "conclusions": "\\label{conc}  We have calculated the signals expected from positron annihilation in galaxy clusters for positrons injected as primaries by embedded radio jets and those produced as secondaries in collisions between cosmic rays and thermal protons in the cluster.  The former case is of interest in constraining the matter content of relativistic jets, while the latter is of interest in measuring the cosmic ray content of the ICM.  The positron annihilation line is particularly interesting for these purposes because typical cluster temperatures $k_B T_e \\sim 1$--$10 \\keV$ are just in the range in which positrons annihilate efficiently without forming positronium.  We calculate the annihilation rates assuming that the cooling positrons mix efficiently with the ambient medium on cosmological timescales.  For positron injection by a single AGN, the peak emissivity in the annihilation line is \\begin{eqnarray} \\dot{n}_{\\rm line,\\, single} & \\approx & 2 \\times 10^{-25} X_s^{\\rm AGN} \\left( \\frac{n_e}{10^{-3} \\cmden} \\right) \\left( \\frac{\\xi_{\\rm AGN}}{0.1} \\right) \\left( \\frac{r_{\\rm mix}}{200 \\kpc} \\right)^{-3} \\nonumber \\\\ \\, & \\, & \\times \\left( \\frac{L_K}{10^{45} \\ergs} \\right) \\left( \\frac{\\tau}{10^8 \\yr} \\right) \\cmdensec, \\qquad \\mbox{\\emph{(single AGN event)}} \\label{eq:agnlineem} \\end{eqnarray} with $X_s^{\\rm AGN} = 1$ $(10)$ for an $s=2$ $(3)$ positron injection spectrum.  The dependence on $s$ appears primarily because, with the normalization procedure described in \\S \\ref{agnsoln}, the total number of $e^+$ produced increases as $s$ increases.  We note that rich clusters can have core densities $n_e \\sim 10^{-2} \\cmden$ and $\\xi_{\\rm AGN}$ may reach a value of several tens of percent for jets with a low bulk Lorentz factor. The time lag between injection and peak annihilation depends principally on the electron density $n_e$ in the cluster core: the characteristic annihilation time of the positron population is $\\tau_{\\rm ann} \\sim 4 \\times 10^9 (n_e/10^{-3}~{\\rm cm^{-3}})^{-1} \\yr$, and the emissivity maintains the peak level for roughly this time period (see Figures \\ref{fig:linetimes2} and \\ref{fig:linetimes3}).  Therefore, although dense clusters produce the strongest signals, they fade relatively quickly.  The longevity of the positron population in typical clusters indicates that the signal can be enhanced if we consider either multiple injection epochs or multiple AGN in a single cluster.  A model of the former type has been recently advocated by \\citet{bohringer} as a way of balancing the cooling flow radiation in some clusters. These authors suggest quasi-periodic mechanical energy injection from the central galaxy throughout the lifetime of the cluster, with injection epochs lasting $\\sim 10^8 \\yr$ occurring every $\\Delta t_{\\rm inj} \\sim 10^9 \\yr$. In the limit in which $\\tau_{\\rm ann} \\gg \\Delta t_{\\rm inj}$, the annihilation emissivity will reach a quasi-steady state with $\\dot{n}_{\\rm line} \\sim \\eta^{-1} \\dot{n}_{\\rm line,\\,single}$, where $\\eta = \\Delta t_{\\rm inj}/\\tau_{\\rm ann}$ is the period between outbursts in units of the annihilation time of the cluster. Here $L_K$ and $\\tau$ are to be interpreted as the mechanical luminosity and lifetime of a single outburst.  An even more interesting possibility is positron injection by galaxies throughout the cluster.  While the vast majority of these galaxies are not active at the present day, recent studies have found evidence for relic supermassive black holes in nearly all bulge-dominated galaxies \\citep{magorrian,gebhardt}.  This, together with modeling of the quasar luminosity function, suggests that nearly all galaxies once hosted a quasar \\citep{haiman,haehnelt}.  Current observations give the relation $M_{\\rm BH} = 7.8 \\times 10^7 (L_{B,\\,{\\rm bulge}}/10^{10} \\lbsun)^{1.08} \\msun$ \\citep{kormendy}.  Because the observed black hole mass-bulge luminosity relation is nearly linear, for the purposes of an estimate it suffices to assume a linear relation and scale our results with the total cluster core luminosity $L_{B,\\,{\\rm cl}}$.  A typical rich cluster core has a total $B$-band luminosity of $L_{B,\\,{\\rm cl}} \\sim 10^{12} \\lbsun$ \\citep{peebles}.  [This luminosity corresponds to $\\sim 100 L_B^{\\star}$ galaxies in the cluster, where $L_B^{\\star}$ is the characteristic luminosity of galaxies in the Schechter function \\citep{yasuda}.] We further assume a mass-to-energy conversion efficiency $\\varepsilon_{\\rm BH}$ during the black hole formation process and that a fraction $f_K$ of this energy is released in outflows.  We expect $\\varepsilon_{\\rm BH} \\sim 0.1$, and observations indicate that $f_K \\sim 0.1$ \\citep[and references therein]{hooper,bigm}. Therefore, if we make the extreme assumption that all injection events occur simultaneously, the peak line emissivity would be \\begin{eqnarray} \\dot{n}_{\\rm line} & \\approx & 9 \\times 10^{-23} X_s^{\\rm AGN} \\left( \\frac{n_e}{10^{-3} \\cmden} \\right) \\left( \\frac{\\xi_{\\rm AGN}}{0.1} \\right) \\left( \\frac{r_{\\rm mix}}{200 \\kpc} \\right)^{-3} \\nonumber \\\\ \\, & \\, & \\times \\left( \\frac{f_K}{0.1} \\, \\frac{\\epsilon_{\\rm BH}}{0.1} \\, \\frac{L_{B,\\,{\\rm cl}}}{10^{12} \\lbsun} \\right) \\cmdensec, \\qquad \\mbox{ \\emph{(multiple simultaneous AGN)}} \\label{eq:multagn} \\end{eqnarray} with $X_s^{\\rm AGN}$ defined as above.  Of course, we must keep in mind that the quasar era peaked at $z \\sim 2$ (\\citealt{pei}; in our cosmology, $t_0 - t_i \\sim 10^{10} \\yr$) so a substantial fraction of the population may have annihilated before the present day. In the opposite limit, in which the source evolution time is much larger than the annihilation time, but in which the annihilation time is in turn much larger than the time between injection events (or $H_0^{-1} \\gg \\tau_{\\rm ann} \\gg \\Delta t_{\\rm inj}$), a quasi-steady state will be reached as described in the previous paragraph, with a steady annihilation emissivity of approximately $\\eta^{-1}$ times that of a typical AGN in the cluster.  For positron production by cosmic rays, the steady-state emissivity in the annihilation line is \\bq \\dot{n}_{\\rm line} \\approx 10^{-30} X_s^{\\rm CR} \\, T_{\\rm keV} \\left( \\frac{n_e}{10^{-3} \\cmden} \\right)^2 \\left( \\frac{\\xi_{\\rm CR}}{0.1} \\right) \\cmdensec, \\qquad \\mbox{\\emph{(cosmic ray secondaries)}} \\label{eq:crlineem} \\eq with $X_s^{\\rm CR} = 1$ $(2.5)$ for an $s=2$ $(s=3)$ proton spectrum.  For an $s=2$ input spectrum, the energy lost to positron annihilation over the age of the universe is $\\sim 10^{-6} U_{\\rm CR} (n_e/10^{-3} \\cmden)$, where $U_{\\rm CR}$ is the cosmic ray energy of the cluster. The emission mechanism is inefficient because $\\la 10^{-3}$ of the cosmic ray energy goes into the positrons and because the characteristic initial energy of the positrons is $\\gg m_e c^2$ (see Figure \\ref{fig:qpp}) so that most of the initial positron energy is lost to cooling radiation before annihilation occurs.  We therefore see that direct positron injection by even a weak AGN ($L_K \\ga 10^{41} \\ergs$) would overwhelm the signal from secondary positrons produced by cosmic ray protons.  The observability of the line is of course the critical question. The flux at earth from a cluster with a positron mixing radius $r_{\\rm mix}$ is \\bq F_c = 8 \\times 10^{-7} \\left( \\frac{ r_{\\rm mix} }{200 \\kpc} \\right)^3 \\left( \\frac{D}{100 \\Mpc} \\right)^{-2} \\left( \\frac{\\dot{n}_{\\rm line}}{10^{-24} \\cmdensec} \\right) \\photflux, \\label{eq:earthflux} \\eq where $D$ is the luminosity distance to the cluster.  In the AGN injection case, the emissivity is $\\propto r_{\\rm mix}^{-3}$, and so the flux is independent of the mixing scale (rather it depends on the total number of positrons injected by the AGN).  For steady production by cosmic ray protons, the emissivity is independent of the volume of the cluster core and the flux scales in proportion to the mixing volume.  Before discussing the prospects for detection of this signal with upcoming instruments, we must consider possible contaminating backgrounds.  First, as described in the previous subsection, the relativistic positrons (and electrons) with $\\gamma \\ga 10^4$ in the cluster generate IC radiation at $511 \\keV$. However, cooling depletes this population rapidly; after a time interval $t - t_i \\sim 1.5 \\times 10^8 \\yr$, the maximum Lorentz factor $\\gamma_{\\rm max} < 10^4$ and the electrons and positrons can no longer produce IC radiation at 511 keV.  However, in the case of steady injection from cosmic ray protons, the IC background originates from the equilibrium distribution itself. Figures \\ref{fig:dndkicn2} and \\ref{fig:dndkicn3} show that in this case the IC background can overwhelm the annihilation line unless the protons have a sufficiently steep ($s \\sim 3$) injection spectrum.  Second, we must also consider the diffuse extragalactic background, for which the flux at $\\epsilon_\\gamma \\approx 511 \\keV$ is $F_{\\rm bkgd} \\sim 2 \\times 10^{-5} \\bkgdflux$ \\citep{watanabe}.  For a cluster of a fixed size, the ratio of the line flux from the cluster to the background flux is independent of cluster distance so long as $z\\ll 1$, \\bq \\frac{F_c}{F_{\\rm bkgd}} \\sim 10^2 \\, T_{\\rm keV}^{-1/2} \\left( \\frac{ r_{\\rm mix}}{200 \\kpc} \\right) \\left( \\frac{\\dot{n}_{\\rm line}}{10^{-24} \\cmdensec} \\right). \\label{eq:fluxratio} \\eq Thus, the expected signal from AGN-injected positrons is well above the diffuse background. However, the signal from secondary positrons generated by cosmic rays will be hidden by the background.  While the prospects for observing the annihilation of secondary positrons in the foreseeable future are small (even if the annihilation line can be observed over the IC background), positrons injected by AGN may soon be detectable with spaceborne instruments. The INTEGRAL satellite\\footnote{ See http://astro.estec.esa.nl/SA-general/Projects/Integral/integral.html}, expected to be launched in October 2002, will have spectral capabilities in the energy range of interest.  The SPI instrument is expected to have a $3 \\sigma$ line sensitivity $\\sim 5.1 \\times 10^{-6} \\photflux$ given an integration time of $10^6 \\sec$, but its poor angular resolution ($2.5\\arcdeg$) may lead to background contamination.  The IBIS instrument, with $12\\arcmin$ resolution, is better suited to cluster detection, but it has a $3\\sigma$ line sensitivity of only $\\sim 2 \\times 10^{-5} \\photflux$ (again for an integration time of $10^6 \\sec$).  An even more powerful search could be conducted with  EXIST\\footnote{ See http://exist.gsfc.nasa.gov }, a proposed all-sky hard X-ray survey mission.  It has an expected $5 \\sigma$ line sensitivity of $\\sim 5 \\times 10^{-6}\\photflux$ in the relevant energy range (assuming an integration time of $10^7 \\sec$, the mean exposure time planned for any point on the sky in the mission), and it has an excellent angular resolution of $5\\arcmin$.  If the positrons are injected by AGN, these sensitivity limits are close to the signal we predict for nearby ($D \\la 100 \\Mpc$) clusters with powerful AGN, multiple injection epochs/galaxies, or steep injection spectra.  Deep exposures with INTEGRAL or statistical analyses taking advantage of the full sky coverage of EXIST may reveal weaker sources as well.  A particularly interesting source is the nearby Virgo cluster, at a distance $\\sim 20 \\Mpc$. In the AGN scenario, the annihilation signal in Virgo is comparable to the detection threshold of all three instruments listed above.  The expected peak flux is \\begin{eqnarray} F_{\\rm Virgo} & \\approx & 1.3 \\times 10^{-6} X_s^{\\rm AGN} \\eta^{-1} \\left( \\frac{n_e}{3 \\times 10^{-3} \\cmden} \\right) \\left( \\frac{\\xi_{\\rm AGN}}{0.1} \\right) \\nonumber \\\\ \\, & \\, & \\times \\left( \\frac{L_K}{10^{44} \\ergs} \\right) \\left(\\frac{\\tau}{10^8 \\yr} \\right) \\photflux, \\label{eq:m87flux} \\end{eqnarray} where we have used fiducial values for the luminosity of M87 estimated by \\citet{bohringer} and for $n_e$ by \\citet{nulsen}.  Here $\\eta^{-1}$ represents the contribution from past AGN phases of cluster galaxies [see the discussion accompanying equation (\\ref{eq:multagn})]; in the best case, it could represent an enhancement of more than an order of magnitude.  If $r_{\\rm mix}$ is large, the high resolution instruments may even be able to map spatial variations in the positron component (and indicate whether the positrons are injected solely by the central galaxy or by a larger number of galaxies during the quasar era).  Continuum $\\gamma$-ray emission from M87, the dominant galaxy in Virgo, may contaminate the annihilation line signal.  To date there have been no observations of Virgo in the relevant energy range.  To estimate the contamination from M87, we assume that observed power-law spectrum in the $2$--$10 \\keV$ range extends to $511 \\keV$; using the recent observations of \\citet{bohr-m87}, we expect M87 to produce a flux in the spectral regime of the annihilation line of $\\sim 4 \\times 10^{-8} T_{\\rm keV}^{1/2} \\photflux$ (the temperature of the core of Virgo varies with radius between $k_B T_e = 1$--$3 \\keV$; see \\citealt{bohr-m87}).  Thus, even if the continuum emission from M87 cannot be removed through the use of a high-angular resolution instrument, the annihilation line should still be visible.  Of course, because M87 is still active, any positrons produced in the current outburst phase have most likely not had sufficient time to cool.  We would therefore expect annihilation line emission only if there is a relic population of positrons either from earlier outbursts of M87 or from the other galaxies in Virgo.  Another interesting source is Centaurus A, the nearest ($D \\approx 3.5 \\Mpc$) bright radio galaxy to the Milky Way.  The radio structure is one of the largest in both apparent and absolute size, covering an area $8\\arcdeg \\times 4\\arcdeg$ on the sky (or $480 \\kpc \\times 240 \\kpc$) and is therefore a very attractive candidate for observations \\citep{israel}.  The expected flux in the positron annihilation line is \\begin{eqnarray} F_{\\rm CenA} & \\approx & 10^{-5} X_s^{\\rm AGN} \\eta^{-1} \\left( \\frac{n_e}{10^{-2} \\cmden} \\right) \\left( \\frac{\\xi_{\\rm AGN}}{0.1} \\right) \\nonumber \\\\ \\, & \\, & \\times \\left( \\frac{L_K}{10^{43} \\ergs} \\right) \\left(\\frac{\\tau}{1.4 \\times 10^8 \\yr} \\right) \\photflux, \\label{eq:cenflux} \\end{eqnarray} where we have scaled $L_K$ to the approximate bolometric luminosity of the Centaurus A nuclear source \\citep{chiaberge} and $\\tau$ to the minimum source age estimated by \\citet{saxton}.  Recent data from the Chandra X-ray Observatory indicates that the gas density around Centaurus A is $n_e \\approx 10^{-2} (r/5 \\kpc)^{-1.33}~{\\rm cm^{-3}}$ at $0.5 \\kpc \\ll r \\la 10 \\kpc$ \\citep{kraft}.  The signal is potentially stronger than that of M87; however, several caveats are in order.  First, the gas near the central source has a temperature $k_B T_e \\sim 0.275\\pm 0.03 \\keV$ \\citep{kraft}, within the regime in which positronium formation begins to play a significant role.  More importantly, the scale of the radio emission suggests that a large fraction of the positrons escape the interstellar medium of Centaurus A and mix with the intragroup medium at $r \\ga 100 \\kpc$.  The inferred virial temperature of the group is only $\\sim 0.07 \\keV$ (based on the observed velocity dispersion of group galaxies; \\citealt{vandenbergh}).  At this temperature, the positronium formation rate is approximately equal to the free annihilation rate \\citep{crannell}. The gas density will also be much lower at large radii (if the observed power-low decline at $r \\sim 10 \\kpc$ continues, then $n_e \\sim 2 \\times 10^{-4} \\cmden$ at $r \\sim 100 \\kpc$).  Therefore it is unclear whether annihilation in this region can produce a strong line.  In addition, because the source is still active, the positrons may not yet have had time to cool or mix sufficiently with the surrounding gas. Finally, because the Centaurus A group contains only $\\sim 30$ galaxies \\citep{vandenbergh}, additional positron enrichment from other AGNs is likely to be small.  In summary, we have shown that, although there are a variety of mechanisms for producing positrons in galaxy clusters, only direct injection by AGNs is likely to produce an observable signal.  Therefore, we argue that a positive detection of positron annihilation lines from clusters suspected of harboring dormant AGNs would be a robust indication that radio jets contain an $e^+e^-$ pair plasma."
    },
    "0201/astro-ph0201063_arXiv.txt": {
        "abstract": "Using {\\it ASCA\\/} data, we find, contrary to other researchers using {\\it ROSAT\\/} data, that the X-ray spectra of the VY~Scl stars TT~Ari and KR~Aur are poorly fit by an absorbed blackbody model but are well fit by an absorbed thermal plasma model. The different conclusions about the nature of the X-ray spectrum of KR~Aur may be due to differences in the accretion rate, since this star was in a high optical state during the {\\it ROSAT\\/} observation, but in an intermediate optical state during the {\\it ASCA\\/} observation. TT~Ari, on the other hand, was in a high optical state during both observations, so directly contradicts the hypothesis that the X-ray spectra of VY~Scl stars in their high optical states are blackbodies. Instead, based on theoretical expectations and the {\\it ASCA\\/}, {\\it Chandra\\/}, and {\\it XMM\\/} spectra of other nonmagnetic cataclysmic variables, we believe that the X-ray spectra of VY~Scl stars in their low and high optical states are due to hot thermal plasma in the boundary layer between the accretion disk and the surface of the white dwarf, and appeal to the acquisition of {\\it Chandra\\/} and {\\it XMM\\/} grating spectra to test this prediction. ",
        "introduction": "Cataclysmic variables (CVs) are a diverse class of semidetached binaries including novae, dwarf novae, and novalike variables, composed typically of a low-mass main-sequence secondary and a white dwarf. With the exception of novae in outburst, the engine for all CVs is the release of gravitational potential energy as material accretes onto the white dwarf. In nonmagnetic systems accretion is mediated by a disk, and simple theory predicts that half of the gravitational potential energy of the accreting material is liberated in the disk and half is liberated in the boundary layer between the disk and the surface of the white dwarf, with luminosities $L_{\\rm disk} \\approx L_{\\rm bl} \\approx G\\Mwd\\Mdot /2\\Rwd =4\\times 10^{34}  (\\Mdot /10^{-8}~\\Msun~{\\rm yr^{-1}})(\\Mwd/\\Msun ) (\\Rwd/10^9~{\\rm cm})^{-1}~\\rm erg~s^{-1}$, where $\\Mdot $ is the mass-accretion rate and $\\Mwd $ and $\\Rwd $ are respectively the mass and radius of the white dwarf. When $\\Mdot $ is low (e.g., dwarf novae in quiescence), the boundary layer is optically thin and quite hot (of order the virial temperature $T_{\\rm vir}=G\\Mwd m_{\\rm H}/6k\\Rwd\\sim 20$ keV); when $\\Mdot $ is high (e.g., novalike variables and dwarf novae in outburst), the boundary layer is optically thick and quite cool (of order the blackbody temperature $T_{\\rm bb}=[G\\Mwd\\Mdot /8\\pi\\sigma\\Rwd ^3]^{1/4} \\sim 10$ eV). In the context of this theory, high-$\\Mdot $ CVs are X-ray sources only because the upper ``atmosphere'' of the boundary layer remains optically thin.  The observational picture of the  X-ray emission of CVs was first sketched in the 1970s and has become clearer through the years as instruments with higher effective area, broader bandpass, and better spectral resolution have flown on X-ray satellites. Summaries of {\\it HEAO~1\\/}, {\\it Einstein\\/}, {\\it EXOSAT\\/}, and {\\it ROSAT\\/} investigations are provided by \\citet{cor81, cor84, era91, muk93, ric96, tee96}; and \\citet{ver97}. That the X-ray emission region of nonmagnetic CVs is compact and centered on the white dwarf is established directly by the X-ray light curves of eclipsing systems. That the X-ray spectra of CVs are due to hot thermal plasma was established by {\\it ASCA\\/}, which had the combination of large effective area, broad bandpass, and high spectral resolution needed to resolve the emission lines of K-shell Mg, Si, S, Ar, and Fe at high energies, and to establish the presence of emission lines of K-shell O and Ne and L-shell Fe at low energies \\citep{nou94, muk00, bas01}. As evidenced by {\\it EUVE\\/} light curves and spectra of the dwarf novae SS~Cyg, U~Gem, VY~Hyi, and OY~Car in outburst, the blackbody component of high-$\\Mdot $ nonmagnetic CVs is typically not observed in the canonical X-ray bandpass (SS~Cyg in outburst is the only clear exception, see \\citealt{beu93, pon95, mau95}) because its temperature is too low and its luminosity is sometimes anomalously weak \\citep[][and references therein]{mau02}. While it is not yet possible to be entirely certain, the working hypothesis is that the X-ray spectra of {\\it all\\/} nonmagnetic CVs are due to hot thermal plasma in the boundary layer between the accretion disk and the surface of the white dwarf. ",
        "conclusions": "Based on an analysis of {\\it ASCA\\/} SIS and GIS spectra, we find that the 0.5--10 keV X-ray spectra of TT~Ari and KR~Aur are poorly fit by an absorbed blackbody model but are well fit by an absorbed thermal plasma model. The TT~Ari spectra are adequately described by a two-temperature solar-abundance thermal plasma, with one component at $kT\\approx 7$~keV and another at $kT\\approx 0.7$ keV, with a relative emission measure of approximately 20:1 in favor of the high-temperature component. The lower-quality KR~Aur spectra are adequately described by a single-temperature solar-abundance thermal plasma with $kT\\approx 6$~keV, and allow a second temperature component only if the column density is significantly increased and the emission measures of the two components are comparable. These results should be understood to be simply {\\it parameterizations\\/} of the X-ray spectra of these two stars, not a definitive determination of the nature of their X-ray spectra. Higher-quality data (e.g., {\\it Chandra\\/} High-Energy Transmission Grating and {\\it XMM\\/} Reflection Grating Spectrometer and European Photon Imaging Camera spectra) are required to determine---by resolving the L-shell emission lines of Fe and the K-shell emission lines of H- and He-like ions of abundant elements from C to Fe---the true emission measure distribution, abundances, density, and Doppler broadening of the plasma in these stars. It is clear, however, that {\\it the X-ray spectra of neither of these stars is that of a blackbody\\/}, contrary to the conclusions of \\citet{sch95} and \\citet{gre98}. The different conclusions about the nature of the X-ray spectrum of KR~Aur may be due to differences in the accretion rate, since this star was in a high optical state during the {\\it ROSAT\\/} observation, but in an intermediate optical state during the {\\it ASCA\\/} observation. TT~Ari, on the other hand, was in a high optical state during both observations, so directly contradicts the claim that the X-ray spectra of VY~Scl stars in their high optical states are blackbodies. It remains possible that TT~Ari, the brightest VY~Scl star, is anomalous, and that the X-ray spectra of other high-state VY~Scl stars are blackbodies, but based on theoretical expectations and the {\\it ASCA\\/}, {\\it Chandra\\/}, and {\\it XMM\\/} X-ray spectra of other high-$\\Mdot $ nonmagnetic CVs, we believe that they are not. Instead, we believe that the X-ray spectra of VY~Scl stars in their low and high optical states are due to hot thermal plasma in the boundary layer between the accretion disk and the surface of the white dwarf, and appeal to the acquisition of {\\it Chandra\\/} and {\\it XMM\\/} grating spectra to test this prediction."
    },
    "0201/astro-ph0201549_arXiv.txt": {
        "abstract": "{ We present the first {\\rm numerical radiative transfer simulation of multiple light scattering} in dust configurations containing aligned non-spherical (spheroidal) dust grains. {\\rm     Such models are especially important if one wants to explain the circular polarization of light, observed in a variety of astronomical objects.} The optical properties of the spheroidal grains are calculated using the method of separation of variables developed by Voshchinnikov \\& Farafonov~(\\cite{vf93}). The radiative transfer problem is solved on the basis of the Monte Carlo method. Test simulations, confirming the correct numerical implementation of the scattering mechanism, are presented. As a first application, we investigate the linear and circular polarization of light coming from a spherical circumstellar shell. This shell contains perfectly aligned prolate or oblate spheroidal grains. We investigate the dependence of the results on the grain parameters (equivolume radius, aspect ratio) and the shell parameters (inner/outer radius, optical thickness). The most remarkable features of the simulated linear polarization maps are so-called polarization null points where the reversal of polarization occurs. They appear in the case when the grain alignment axis is perpendicular to the line of sight. {\\rm     The position of these points may be used for the estimation of grain shape and geometrical structure of the shell. The origin of null points lies in the physics of light scattering by non-spherical particles and is not related to the cancellation of polarization as  was discussed in previous models.} The maps of circular polarization have a sector-like structure with maxima at the ends of lines inclined to the grain alignment axis by $\\pm 45\\degr$. ",
        "introduction": "It is now well established that the polarization of optical {\\rm and near-infrared radiation} from young and evolved stellar objects, reflection nebulae, and active galactic nuclei is mainly caused by dust grains. In many cases the observed polarization can be satisfactorily interpreted by light scattering on spherical grains. In particular, the orientation of the polarization vectors is often used for {\\rm the investigation of the dust distribution and for identifying the location of the embedded illuminating source(s).}  However, even in the case of  simple objects like ordinary reflection nebulae, deviations of polarization vectors from the direction perpendicular to a star were discovered more than three decades ago (Elvius \\& Hall~\\cite{eh67}). These deviations are displayed in the outer filamentary parts of the Merope nebula {\\rm and may be easily explained by light scattering on non-spherical grains aligned by a magnetic field  along the filaments.} Non-centrosymmetric polarization patterns have been observed in bipolar and cometary nebulae (Scarrott et al.~\\cite{sdw89}), young stellar objects (Hajjar \\& Bastien~\\cite{hb96}), evolved stars (Kastner \\& Weintraub~\\cite{kw96}), and are also clearly seen in polarization maps of comets (Dollfus \\& Suchail~\\cite{ds87}). {\\rm They may be attributed to non-spherical grains, although they could be caused by multiple scattering on spherical particles as well.} Another effect which may be related to the light scattering by non-spherical grains is the wavelength dependence of the positional angle of polarization observed in red giants, AGB stars, and bipolar  reflection % nebulae (see, e.g., Johnson \\& Jones~\\cite{jj91}). The variations from blue to red may reach $20\\degr - 60\\degr$ and it is very difficult or even impossible to interpret this behaviour using spherical grains only.  Recently, very high degrees of circular polarization of scattered light in the Orion molecular cloud were measured by Chrysostomou et al.~(\\cite{cgm00}). The authors suggest that the circular polarization is produced by aligned non-spherical grains. {\\rm Multiple scattering of radiation by spherical or randomly oriented non-spherical grains results in a much smaller circular polarization degree than observed.} In addition, the interstellar polarization and polarized thermal emission phenomena prove that non-spherical grains exist in the interstellar medium. These effects arise because of dichroic extinction/emission of radiation by aligned non-spherical grains and were modeled with spheroidal grains (Voshchinnikov~\\cite{v90}; Kim \\& Martin~\\cite{km95}; Onaka~\\cite{o00}).  Up to now, spheroidal grains have been used for the interpretation of scattered radiation only in the case of single scattering (Voshchinnikov~\\cite{v98}; Gledhill \\& McCall~\\cite{gm00}). {\\rm The numerical simulation of the polarized radiation transfer through a medium with non-spherical grains, including multiple scattering, is extremely difficult. In the case of spherical grains nearly all previous simulations are based on Monte Carlo simulations.} This method was used for the interpretation of polarimetric observations of various stars (see, e.g., Daniel~\\cite{d80}; Voshchinnikov \\& Karjukin~\\cite{vk94}) as well as the production of polarimetric maps of different extended objects (e.g., Bastien \\& M\\'enard~\\cite{bm88}; Fischer et al.~\\cite{fhy94}; Wolf et al.~\\cite{wfp98}). The most recent, major achievement of the Monte Carlo technique --- in respect to the solution to the radiative transfer (RT) problem --- is the development of RT codes that allow to calculate the spectral energy distribution of multi-dimensional dust configurations self-consistently (Wolf et al.~\\cite{whs99}; Wolf \\& Henning~\\cite{wh00}).  {\\rm The main aim of this paper is to provide the first radiative transfer simulations including scattering, extinction,  absorption, and re-emission of radiation by aligned spheroidal dust grains. Our formalism will be applied to several simple model configurations. Because of the complexity of the problem, we decided to consider the separate mechanisms step by step. In this paper,  the basic theory is presented and the main numerical features of the simulation are described (Sect.~\\ref{basic-all}). Furthermore, we consider multiple scattering of light by spheroidal grains and the resulting linear and circular polarization (Sect.~\\ref{test} and \\ref{appl}).} ",
        "conclusions": "The paper contains the first solution to the RT problem including multiple light scattering in dust configurations with aligned non-spherical dust grains. Our main aim was  to discuss the effects of light scattering by spheroidal dust grains in a simple (spherical) dust configuration, containing perfectly aligned particles.  The most remarkable features of the simulated linear polarization maps are the polarization null points where the reversal of polarization occurs. They appear when the grain alignment axis is perpendicular to the line of sight. Symmetrically to the polarization null points, we found maxima in the intensity maps.  In contrast to spherical grains, even single scattering by spheroidal grains may cause circular polarization. The maps of circular polarization have a sector-like structure with polarization maxima at the ends of a line inclined to the grain alignment axis by $\\pm 45\\degr$.  {\\rm Based on our theoretical investigations of light scattering, dichroic absorption, and (re)emission by spheroidal grains presented in Sect.~\\ref{basic-all}, the next steps of the numerical radiative transfer simulations will include the consideration of light scattering by partly aligned rotating particles and the calculation of the polarized re-emission of grains in different dust configurations. This will immediately lead to an improved interpretation and therefore better understanding of mid-infrared and submillimeter polarization % where re-emission and dichroic absorption by partly aligned non-spherical dust grains are the main polarization mechanisms. Thus, it provides a basis for investigations of the structure of the magnetic fields in these objects which are in most cases assumed to cause grain alignment. }"
    },
    "0201/astro-ph0201255_arXiv.txt": {
        "abstract": "\\noindent Estimation of distances to nearby galaxies by the use of eclipsing binaries as standard candles has recently become feasible because of new large scale instruments and the discovery of thousands of eclipsing binaries as spinoff from Galactic microlensing surveys. Published measurements of distances to detached eclipsing binaries in the Large Magellanic Cloud combine stellar surface areas (in absolute units) determined from photometric light and radial velocity curves with surface brightnesses from model atmospheres and observed spectra. The method does not require the stars to be normal or undistorted, and is not limited in its applicability to the well detached systems that have traditionally been considered. We discuss the potential usefulness of semi-detached vis \\`{a} vis detached eclipsing binaries for distance determination, and examine and quantify criteria for their selection from large catalogs. Following our earlier paper on detached binaries in the Small Magellanic Cloud (SMC), we carry out semi-detached light curve solutions for SMC binaries discovered by the OGLE collaboration, identify candidates for SMC distance estimation that can be targets of future high quality observations, and tabulate results of OGLE light curve solutions. We point out that semi-detached binaries have important advantages over well-detached systems as standard candles, although this idea runs counter to the usual view that the latter are optimal distance indicators. Potential advantages are that (1) light curve solutions can be strengthened by exploiting lobe-filling configurations, (2) only single-lined spectra may be needed for radial velocities because the mass ratio can be determined from photometry in the case of complete eclipses, and (3) nearly all semi-detached binaries have sensibly circular orbits, which is not true for detached binaries. We carry out simulations with synthetic data to see if semi-detached binaries can be reliably identified and to quantify the accuracy of solutions. The simulations were done for detached as well as semi-detached binaries so as to constitute a proper controlled study. The simulations demonstrate two additional advantages for semi-detached distance determination candidates; (4) the well-known difficulty in distinguishing solutions with interchanged radii (aliasing) is much less severe for semi-detached than for detached binaries, and (5) the condition of complete eclipse (which removes a near degeneracy between inclination and the ratio of the radii) is identified with improved reliability. In many cases we find that parameters are accurately determined (\\textit{e.g.} relative errors in radii smaller than 10\\%), and that detached and semi-detached systems can be distinguished. We select 36 candidate semi-detached systems (although 7 of these are doubtful due to large mass ratios or periods) from the OGLE SMC eclipsing binary catalog. We expected that most semi-detached candidates would have light curves similar to those of common Algol binaries but that turned out not to be the case, and we note that fully Algol-like light curves are nearly absent in the OGLE sample. We discuss possible explanations for the near absence of obvious Algols in OGLE, including whether their paucity is real or apparent. ",
        "introduction": "\\noindent  Accurate measurement of the distances to the Magellanic Clouds is an important current issue as it provides a basic step toward determining the extragalactic distance scale. Paczynski (1997, 2000) has argued that eclipsing binaries now provide the most direct and accurate distances to the Magellanic Clouds, and examples are already in the literature (e.g. Guinan et al. 1998; Fitzpatrick et al. 2001). Recently, thousands of variable stars including eclipsing binaries have been discovered by the OGLE (Udalski et al. 1998), MACHO (Alcock et al. 1997) and EROS (Grison et al. 1995) collaborations as a by-product of galactic microlensing searches. These catalogs motivate a systematic, quantitative search for close to ideal systems for distance determination.  The method of measuring distances by means of eclipsing binaries has been known for decades, and its basis has been clearly explained by Paczynski (1997) and Guinan, et al. (1998), among others. In essence, light curves provide relative star dimensions ($R_{1}/a$, $R_{2}/a$, where $R$ is mean radius and $a$ is orbit size) and radial velocities establish the absolute scale by providing the orbit size, so that one can find the $R$'s in physical units by combining the two kinds of information. Fine effects, such as departures from sphericity, etc., can be modeled by modern eclipsing binary light curve programs. With absolute radii known, luminosities in physical units follow if emission per unit surface area (energy per unit area per unit time per wavelength interval) becomes known. Some persons favor calibrated relations based on interferometrically resolved, un-complicated stars for the emission measure while others favor the predictions of stellar atmosphere models that are fitted to spectral energy distributions (SED) of eclipsing binary distance estimation targets. Emission for plane-parallel atmospheres is determined, in principle, if effective temperature, $\\log$ $g$, and chemical composition are specified. For a well observed eclipsing binary, $g$ is computed to better than adequate accuracy as $GM/R^{2}$ and $T_{eff}$ can be estimated by fitting a theoretical SED to an observed SED, as in Fitzpatrick \\& Massa (1999). In usual practice, surface chemical composition would be assumed normal, but that is the \\textit{only} normalcy assumption. The reasonableness of that assumption for \\emph{SD} binaries will be discussed below. The overall method is mainly geometrical, with only the emission measure involving radiative physics. It does not require knowledge of distances to calibration stars, as opposed to other standard candle methods such as by Cepheid variables or supernovae, and therein lies one of the primary advantages. Empirical calibration errors are bypassed if surface emission is computed from a stellar atmosphere model.  Although conventional wisdom holds that well-detached eclipsing binaries yield the most reliable light curve solutions, the basis for that conjecture may not extend beyond the scientific instinct that simpler is better. In fact there are real advantages to solutions of semi-detached (hereafter \\emph{SD}) and overcontact (\\emph{OC}) binary light curves, partly in the exploitation of lobe-filling configurations and partly through proximity effects, which provide information that is lacking in well-detached binaries. Actually, many factors influence the relative reliability of detached, \\emph{SD}, and \\emph{OC} light curve solutions. Accordingly, searches for standard candle binaries should examine all relevant considerations, including ones that argue for or against \\emph{SD} and \\emph{OC} systems. Here we consider \\emph{SD} binaries in the Small Magellanic Cloud (SMC) and will take up the \\emph{OC} case in a forthcoming paper. Our aims are to discuss the main considerations that bear upon the potential usefulness of \\emph{SD} binaries as standard candles, to identify good \\emph{SD} candidates for SMC distance determination via future observations with large telescopes, and to derive preliminary dimensional, radiative, and mass ratio properties of the candidates. Of course, \\emph{SD} binaries are fascinating objects in their own right, and we expect that their identification will also lead to investigations of \\emph{SD} properties unrelated to distance determination. SMC detached binaries were treated in Wyithe \\& Wilson~(2001, hereafter Paper I).  The remainder of the paper is in four parts. Section 2 considers potential advantages of \\emph{SD} binaries as standard candles. Sec.~\\ref{fit_scheme} discusses our automated fitting scheme and differences from the scheme described in Paper I for detached binaries. Quite apart from the logical arguments of Section 2, simulations can show statistically how well \\emph{SD} and detached solutions recover known parameters. This topic is discussed in Sec.~\\ref{mocksec}. We also show how in some cases the issue of whether a binary is detached or \\emph{SD} can be determined from photometry with reasonable reliability. Sec.~\\ref{OGLEfits} has solutions to the OGLE catalog of eclipsing binary stars and discusses candidate \\emph{SD} binaries. ",
        "conclusions": "Although traditionally used, well detached eclipsing binaries are not necessarily the ideal or only choice for eclipsing binary distance determination, since absolute brightnesses on absolute stellar surface elements can be computed and integrated over surfaces for almost all classes of eclipsing binaries. In this paper we have presented arguments in support of semi-detached (\\emph{SD}) eclipsing binaries as standard candles, and taken this as motivation to find \\emph{SD} solutions to the OGLE SMC eclipsing binary catalog, and to select \\emph{SD} systems for future study.   Several advantages of \\emph{SD} binaries, in particular the exploitation of lobe-filling configurations lead to accurate light-curve solutions and may therefore lead to accurate distances. However before investing the considerable effort to make the observations required for a distance determination, it is helpful to have confidence that the binary is appropriate for the purpose. We have found several advantages for \\emph{SD} systems that relate to the selection of candidates from the large catalogs of light-curves now becoming available. We find that aliased solutions (where the radii are interchanged) are significantly less of a problem for \\emph{SD} than detached systems. Furthermore, the inclination, and therefore the ratios of radii and luminosity, as well as whether the system undergoes complete eclipse are much better determined. Candidate \\emph{SD} distance determination systems can therefore be selected for the desirable properties of having double-lined spectra and complete eclipses more reliably than can detached systems.  We have computed both \\emph{SD} and detached solutions to the 1459 eclipsing binary stars identified in the SMC by the OGLE collaboration (Udalski et al.~1998). This work follows our earlier paper on detached systems. By fitting simulated catalogs we estimate a success rate of 98 percent for finding acceptable converged solutions to \\emph{SD} configurations. Acceptance of an \\emph{SD} solution does not establish that a system surely is \\emph{SD}, as detached systems can mimic \\emph{SD} light-curves and vice-versa. However, we show that the system condition (detached or \\emph{SD}) can often (for about 1/3 of our simulated binaries) be determined from the ratio of residuals (\\textit{SS} ratio). Of the OGLE systems with both kinds of solution, 36 percent have \\textit{SS} ratios significantly different from unity, allowing morphological categorization. In particular, also requiring eccentricity consistent with zero, we find 36 systems that can be identified as being of the \\emph{SD} morphological type with reasonable reliability (although 7 of these are doubtful due to very large mass ratios or periods), such that future observations with large scale optics should lead to accurate distance determinations. As emphasized in Section 4, we anticipate a coming time when several tests will better sift through high quality light curves and velocities from large optics. Comparison of photometric and spectroscopic mass ratios will then settle most of the otherwise unclear decisions between detached and \\emph{SD} assignments.   Although we expected that most binaries selected as \\emph{SD} would be common normal Algols, the result is that only a small minority of our \\emph{SD}'s have light curves like those of common Algols. Indeed, inspection of the OGLE catalog reveals that systems with light-curves resembling those of classical Algols are virtually absent. Only two of our \\emph{SD} candidate binaries have solutions consistent with normal Algols. Basically the secondary stars have high surface brightnesses (\\textit{i.e.} are hot compared to those of normal Algols). This outcome could be a strongly positive one, as \\emph{SD}'s with bright secondaries will have stronger light curve solutions (with information from \\textit{two} deep eclipses) than ordinary Algols, while they may retain all the \\emph{SD} advantages mentioned in the Introduction. They also have moderate luminosity ratios, and therefore greatly increased likelihood of their spectra being double-lined.  It could be that some of the 36 \\emph{SD} binaries actually are \\emph{OC}, as testing for the \\emph{OC} condition is beyond the scope of this paper and will be treated in future work. That outcome also could be advantageous, as \\emph{OC} configurations can have very well conditioned light curve solutions, perhaps even better than \\emph{SD}'s. \\emph{OC} binaries are rare at the high luminosities detectable by OGLE (although very common at much lower luminosity), yet discovery of just a few would be valuable help in finding the distance to the Magellanic Clouds.  The next step is to obtain spectra and accurate multi-band light curves of the more promising systems so as to confirm \\emph{SD} assignments via spectral types and improved light curve parameters, and also to observe radial velocities for absolute dimensions."
    },
    "0201/hep-th0201042_arXiv.txt": {
        "abstract": "We propose that higher-dimensional extended objects ($p$-branes) are created by super-Planckian scattering processes in theories with TeV scale gravity. As an example, we compute the cross section for $p$-brane creation in a $(n+4)$-dimensional spacetime with asymmetric compactification. We find that the cross section for the formation of a brane which is wounded on a compact submanifold of size of the fundamental gravitational scale is larger than the cross section for the creation of a spherically symmetric black hole. Therefore, we predict that branes are more likely to be created than black holes in super-Planckian scattering processes in these manifolds. The higher rate of $p$-brane production has important phenomenological consequences, as it significantly enhances possible detection of non-perturbative gravitational events in future hadron colliders and cosmic rays detectors. \\\\\\\\ PACS: 04.80.Cc; 04.50.+h; 11.27.+d; 11.80.-m; 13.85.Tp; 13.85.-t ",
        "introduction": " ",
        "conclusions": ""
    },
    "0201/astro-ph0201019_arXiv.txt": {
        "abstract": "We develop a detailed chemical model for the starless cores of strongly magnetized molecular clouds, with the ambipolar diffusion-driven dynamic evolution of the clouds coupled to the chemistry through ion abundances. We concentrate on two representative model clouds in this initial study, one with magnetic fields and the other without. The model predictions on the peak values and spatial distributions of the column densities of CO, CCS, N$_2$H$^+$ and HCO$^+$ are compared with those observationally inferred for the well-studied starless core L1544, which is thought to be on the verge of star formation. We find that the magnetic model, in which the cloud is magnetically supported for several million years before collapsing dynamically, provides a reasonable overall fit to the available data on L1544; the fit is significantly worse for the non-magnetic model, in which the cloud collapses promptly. The observed large peak column density for N$_2$H$^+$ and clear central depression for CCS favor the magnetically-retarded collapse over the free-fall collapse. A relatively high abundance of CCS is found in the magnetic model, resulting most likely from an interplay of depletion and late-time hydrocarbon chemistry enhanced by CO depletion. These initial results lend some support to the standard picture of dense core formation in strongly magnetized clouds through ambipolar diffusion. They are at variance with those of Aikawa et al. (2001) who considered a set of models somewhat different from ours and preferred one in which the cloud collapses more or less freely for L1544. ",
        "introduction": "Through a close interplay between theory and observation, a ``standard'' picture for the formation of isolated, low-mass stars has emerged (Shu, Adams \\& Lizano 1987). At the heart of this picture lie the so-called ``dense cores'', studied extensively by Myers and coworkers (e.g., Myers 1999). These dense cores are intimately associated with star formation, with roughly half of them already harboring infrared sources (Beichman et al. 1986). The other half are termed ``starless cores'', and they are the focus of our investigation.  The starless cores are thought to be condensed out of strongly magnetized, turbulent background clouds. Direct Zeeman measurements to date, as compiled by Crutcher (1999), suggest that the field strength is close to the critical value required for the molecular cloud support, after likely geometric corrections (Shu et al. 1999). Core formation can be driven by either turbulence decay (Myers 1999) or ambipolar diffusion. We shall focus on the latter, which has been studied quantitatively by many authors (e.g., Nakano 1979; Lizano \\& Shu 1989; Basu \\& Mouschovias 1994; Ciolek \\& Mouschovias 1994). A general conclusion is that dense cores are formed on a time scale several times the free fall time of the background. This relatively long formation time should leave a strong imprint on the chemistry of the cores.  The goal of our investigation is to predict the spatial distributions of various molecular species at different stages of cloud evolution. For this initial study, we will concentrate on CO, CS, CCS, NH$_3$, N$_2$H$^+$ and HCO$^+$, which are often used to probe the physical conditions of star-forming clouds. We take into account the feedback of chemistry on the cloud dynamics, through ion abundances, which regulate the ambipolar diffusion and thus cloud contraction rate. Through a detailed comparison of the predicted and observed chemical abundances of starless cores, we seek to provide further support for the standard picture of isolated star formation involving ambipolar diffusion. A starless core well suited for such a purpose is the dense core of L1544, for which an extensive data set is becoming available (e.g., Caselli et al. 2001a,b).  The L1544 starless core is located in the eastern part of the Taurus molecular cloud complex, at an estimated distance of 140 pc (Elias 1978). It is an elongated core that shows substantial infall motions (up to $\\sim 0.1$ km s$^{-1}$) on both large ($\\sim 0.1$ pc) and small ($\\sim 0.01$ pc) scales, based on single-disk observations of CO, $^{13}$CO, C$^{18}$O, CS, C$^{34}$S, HCO$^+$, and N$_2$H$^+$ (Tafalla et al. 1998) and interferometric observations of N$_2$H$^+$ (Williams et al. 1999). Interferometric observations of CCS by Ohashi et al. (1999) reveal a ring-like structure, which shows evidence for collapse as well as rotation. The low central emission is most likely due to a depletion of CCS, judging from the fact that dust continuum peaks inside the ring (Ward-Thompson, Motte \\& Andre 1999). Indeed, the column density distribution of the L1544 core is rather well determined\\footnote{See, however, Evans et al. (2001) who find that a singular isothermal sphere cannot be ruled out from the modeling of SCUBA data alone, when the lower temperature at the core center relative to edge is taken into account.}, not only from the dust continuum emission, but also from absorption against a mid-infrared background (Bacmann et al. 2000). The core has a small, high column (and volume) density central plateau (of radius $\\sim 2900$ AU), surrounded by an envelope in which the column (and volume) density decreases rapidly outward, before it joins rather abruptly with the background at a radius of order $10^4$ AU (Bacmann et al. 2000). The plateau-envelope structure has been interpreted in terms of the ambipolar diffusion-driven dynamical evolution of a magnetized cloud (Li 1999; Ciolek \\& Basu 2000), and the large column density contrast between the central plateau and the background ($\\sim 20$; Bacmann et al. 2000) suggests that the L1544 core is rather evolved (Williams et al. 1999), and may be on the verge of forming a star (or stellar system). Further evidence for the core being in an advanced stage of evolution comes from IRAM 30m observations of Caselli et al. (1999), who found that CO is depleted by a factor of order $\\sim 10$ at the dust continuum peak. These authors also noted that the (likely) optically thin lines of D$^{13}$CO$^+$ and HC$^{18}$O$^+$ are double-peaked, which may be another indication that these species are centrally depleted and their emission comes mainly from an infalling {\\it shell}. The wealth of molecular line and dust continuum data, coupled with the recent Zeeman measurement of magnetic field strength (Crutcher \\& Troland 2000) and the determination of magnetic field direction from submillimeter polarization measurements (Ward-Thompson et al. 2000), makes L1544 a Rosette Stone for theories of dense core condensation leading to isolated low-mass star formation.  Many of the dynamical and chemical features of L1544 are shared by L1498, another quiescent, elongated starless core also located in the Taurus molecular complex. Lemme et al. (1995) studied the core in detail, and found that the distributions of both C$^{18}$O and CS are ring-like, indicating a central depletion of these two species. The CO depletion factor is inferred to be of order $\\sim 8$ or higher (Willacy, Langer \\& Velusamy 1998), comparable to that of L1544, based on ISOPHOT 100 and 200 $\\mu$m observations. Single-dish and interferometric observations of CCS (Kuiper, Langer \\& Velusamy 1996) show that its distribution is also ring-like. Moreover, the NH$_3$ emission appears to be centrally peaked, again similar to L1544. Perhaps most strikingly, there is evidence for infall motions on a large ($\\sim 0.1$ pc) scale in the double-peaked profiles of both optically thick CS lines and the likely optically thin C$^{34}$S(2-1) transition (Lemme et al. 1995); the latter could come from an infalling shell, as in L1544 (Caselli et al. 1999). The optically thin tracer N$_2$H$^+$ shows, however, a single peaked spectrum (Lee, Myers \\& Tafalla 2001), indicating that it is centrally peaked, in contrast to C$^{34}$S. These similarities to L1544 point to a relatively advanced stage of pre-protostellar evolution for L1498, making it another valuable testing ground for chemical and dynamical models of (low-mass) star formation.  Common to both sources is the ring (or shell)-like structure of CCS, a well-known ``early time'' species (Bergin \\& Langer 1997) tracing chemically ``young'' material of order $10^5$ yrs or less. This prominent feature should provide a strong constraint on models. Additional constraints should come from the inferred large CO depletion near the core center as well as the spatial extent and speed of the infall motions present in both L1544 and L1498. The question we want to address is: can the standard picture of dense core formation involving ambipolar diffusion satisfy these chemical {\\em and} kinematic constraints simultaneously? A partial answer to this question has been provided by Bergin \\& Langer (1997), who considered the chemical evolution of {\\it a parcel} of cloud material whose density increases with time in a prescribed way (i.e., a standard time dependent but depth independent model). They found that differential depletion of various molecular species onto dust grains must be a key ingredient in understanding the differentiated chemical structure observed in L1498.  Aikawa et al. (2001) went one step further, and followed the evolution of molecular abundances of a dynamically collapsing starless core {\\it as a whole}, with an emphasis on the spatial distribution. The study was based on the self-similar solution of Larson (1969) and Penston (1969) and its artificially delayed analogs. Among the dynamical models adopted, they found that the undelayed Larson-Penston model provides the best overall fit to the observational data on CO, CCS and N$_2$H$^+$ for L1544. If this result is robust, its implications would be far reaching. It is well known that the Larson-Penston solution describes the collapse of a cloud essentially on a free-fall time scale. The fact that it matches observations better than its delayed analogs appears to pose a serious challenge to the standard picture of star formation involving strongly magnetized cloud cores, which are formed through ambipolar diffusion on a time scale much longer than the free-fall time scale. However, the large infall velocity of the Larson-Penston solution, approaching 3.3 times the sound speed, is clearly incompatible with those inferred for L1544 and other starless cores, which are typically less than the sound speed (Lee et al. 2001). Furthermore, the best-fit model of Aikawa et al. under predicts the abundance of N$_2$H$^+$ by a large factor of $\\sim 20$. These discrepancies motivate us to reexamine the molecular evolution of starless cores using a more sophisticated dynamical model, one that is coupled to the cloud chemistry.  The dynamical model we will adopt is that of Li (1999). It is a spherical model of starless cores that takes into account the dynamic effects of magnetic field approximately. As envisioned in standard scenario of low-mass star formation, the cores in the model evolve through ambipolar diffusion, and their dynamics are coupled to chemistry through the abundances of charged species. Previously, we (Shematovich, Shustov, \\& Wiebe 1997; 1999) have modeled the chemistry of dynamically evolving clouds in detail. The chemical model will be combined with the dynamical model into a coupled dynamical-chemical model of magnetized starless cores. An advantage of the combined model is that it allows for a determination of not only the time evolution of various molecular species but also their {\\em spatial} distributions {\\em and} kinematics. We describe the formulation of the combined model in \\S\\ref{review}. In \\S\\ref{result}, we present representative model results, focusing on the density distribution, velocity field and abundance distributions of several commonly used molecular species. These results are compared with observations of L1544. We conclude and discuss our main results and future refinements in \\S\\ref{discuss}. ",
        "conclusions": "\\label{discuss}  \\subsection{Magnetic Cloud Support and CCS Abundance} \\label{ccs}  The success of the standard, magnetic model in fitting the data on L1544 hinges to a large extent on the relatively high abundance predicted for CCS. This result is somewhat surprising in view of the fact that CCS is a well-known ``early time'' species, whose abundance should decline rapidly after a few times $10^5$ years according to Suzuki et al. (1992), Bergin \\& Langer (1997), and Aikawa et al. 2001). Part of the reason for the persistently high CCS abundance in the standard model even after some 5.69~Myrs of evolution (the time it takes the cloud to reach the observed state) is that we adopted an initial cloud density of $n_{\\rm H_2}=10^3$~cm$^{-3}$, which is lower than those of Bergin \\& Langer ($10^{3.5}$~cm$^{-3}$) or Suzuki et al. and Aikawa et al. ($10^4$~cm$^{-3}$). As a result, the bulk of the $5.69$~Myrs is spent in the ``pre-dense core'' phase when the central density remains below $10^4$~cm$^{-3}$ (see Table~3). Only a small fraction ($\\sim 10\\%$) of that time is spent at densities above $10^4$~cm$^{-3}$. That a cloud spends a relatively short time (a few times the free-fall time) at high densities after forming a magnetically supercritical core is a general feature of ambipolar diffusion-driven evolution. It helps to keep the CCS abundance higher than one would expect based on the total time of cloud evolution.  The fractional abundances of CCS are shown in Fig.~6 as a function of radius for both the magnetic and non-magnetic models. Note the ``humps'' on the distributions of CCS abundance in the magnetic model, which are not apparent in the non-magnetic model. The exact origin of the humps is unclear. We suspect that they are related to the depletion of CO in the central high density region, similar to the ``depletion'' peak proposed by Ruffle et al. (1997) for HC$_3$N. Ruffle et al. (1997, 1999) showed that C$_2$H and HC$_3$N in dense, cold cores are characterized by a secondary ``late-time'' maximum in their fractional abundances. The second peak is thought to be mainly caused by the increased depletion of CO, which allows C$^+$ ion to react more readily with H$_2$ than with oxygen bearing species. The net result is an increase in the rates of production of CH and other carbon-bearing molecules (without oxygen). It was found that high late-time C$_2$H and HC$_3$N maxima are achieved only when the freeze-out time scale is long compared to the chemical time scale. Because C$_2$H and other (neutral and ionized) late-time hydrocarbons are the precursors of CCS formation in the reactions with neutral and ionized sulfur through \\begin{equation} {\\rm {C_2H,C_2H_2,C_2H_3, ...} + S^{+} \\rightarrow HCCS^{+} + ...., } \\end{equation} \\begin{equation} {\\rm {C_2H_2^{+},C_2H_3^{+},C_2H_4^{+}, ...} + S \\rightarrow HCCS^{+} + ....,} \\end{equation} \\begin{equation} {\\rm HCCS^{+} + e \\rightarrow CCS + H, } \\end{equation} CCS should behave in a way similar to the C$_2$H and HC$_3$N. The competition between depletion and late-time hydrocarbon chemistry enhanced by CO-depletion appears to be the most likely cause of the humps on the radial profiles of CCS abundance.  In any case, the humps have apparently kept the CCS abundance of the magnetic model close to, and in some regions exceeding, that of the non-magnetic model. They are largely responsible for the higher-than-expected value of the CCS column density in the magnetic model. In addition, the decline of CCS abundance outside the hump, coupled with a steep decrease in the hydrogen number density with radius, may explain the observed rapid decline of the CCS column density outside the ring (Ohashi 2001; priv. comm.).  \\subsection{Depletion of Molecules onto Dust Grains}  \\subsubsection{Depletion of CO and Other Species}  In the standard model, CO is heavily depleted in the central high density region. The depletion is most clearly seen in the first panel of Fig.~7, where the fractional abundance is plotted as a function of radius. The abundance is significantly below the canonical value of $4.7\\times 10^{-5}$ inside a radius of $\\sim 10^4$~AU, by a factor up to $\\sim 30$. The size of the depletion region is roughly consistent with that deduced by Caselli et al. (2001b; their Fig.~2). The depletion is less evident in the CO column density distribution (see Fig.~4), which is flat near the center. The ratio of CO and hydrogen in peak column density is a factor of $\\sim 5$ below the canonical value. This depletion factor is about half of the value ($\\sim 9$) inferred by Caselli et al. (2001b). Their inference is based on the observations of C$^{17}$O, C$^{18}$O, and dust emission (Ward-Thompson et al. 1999). In particular, an dust opacity of 0.005~cm$^2$g$^{-1}$ was adopted to convert the dust continuum flux into a hydrogen column density. At densities as high as $10^6$~cm$^{-3}$, an opacity of $\\sim 0.01$~cm$^2$g$^{-1}$ may be more appropriate (Ossenkopf \\& Henning 1994), which would lower the estimate of the hydrogen column density by half and bring a closer agreement between the model prediction and observation. Alternatively, a flattened geometry, as observed for L1544, and/or a somewhat higher CO adsorption energy (see Aikawa et al. 2001) may enhance the CO depletion in the column density distribution.  Besides CO, there are other species that are strongly depleted in the high density central region. These include CS, CCS and to a lesser extent HCO$^+$, as can be seen from Fig.~7. The nitrogen bearing species, N$_2$H$^+$ and NH$_3$, are on the other hand hardly depleted; if anything, their abundances are slightly enhanced near the center. This differential depletion pattern is in agreement with previous findings (e.g., Bergin \\& Langer 1997; Aikawa et al. 2001). It may have profound effects on the molecular line profiles used to probe cloud kinematics.  \\subsubsection{Effects of Depletion on Line Profiles}  An interesting observational fact about L1544 is that some optically thin lines in L1544, such as CCS (Ohashi et al. 1999) and HC$^{18}$O$^+$ (Caselli et al. 2001b), are double-peaked. As noted by Caselli et al. (1999, 2001b), the double-peaked optically thin lines can be explained in a collapsing cloud {\\it provided} that the line-emitting molecules are strongly depleted near the center. We have seen from Fig.~7 that CCS and, to a lesser extent, HCO$^+$ (and thus HC$^{18}$O$^+$) are indeed centrally depleted in our standard model. A unique strength of the coupled dynamical-chemical model is that it allows for a simultaneous determination of the spatial distribution of a given species and the velocity field, the two ingredients for line profile modeling. In a subsequent paper, we will produce synthetic line profiles for both optically thin and thick lines (such as CS and HCO$^+$; Tafalla et al. 1998), which should enable us to constrain the models of core formation leading to star formation using both the spatial {\\it and} kinematic information of molecular species.  \\subsection{Conclusions}  We have developed a coupled dynamical and chemical model for the starless cores of strongly magnetized molecular clouds. The coupling is achieved through ions, the abundances of which are determined self-consistently from a chemical network. The ionic abundances control the time scale of cloud dynamic evolution via ambipolar diffusion. We have concentrated on a representative model in which an isolated magnetic cloud increases its central number density by a factor of $10^3$ to $n_{\\rm H_2}=10^6$~cm$^{-3}$ in 5.69 million years, with an eye on explaining the observational data on the chemical abundances of the well-studied starless core L1544. We find that the predicted peak values of CO, CCS, N$_2$H$^+$ and HCO$^+$ column densities are within a factor of two or so of those inferred for L1544 from observations. The spatial distributions of these species are also consistent with those observed. An alternative, non-magnetic model was also considered for comparison. With the same initial conditions as in the magnetic model, the non-magnetic cloud collapses promptly, reaching the observed state in merely 1.56 million years. It produces a worse overall fit to the available data on L1544. In particular, the column density distribution of CCS predicted by the model is centrally peaked, which is not observed.  We conclude that our initial results of modeling lend some support to the standard picture of dense core formation out of strongly magnetized molecular clouds involving ambipolar diffusion over several dynamic times. There are, however, a number of uncertainties to which the results may be sensitive to, including the sticking probability and adsorption energies for various species. A parameter survey is needed to firm up the conclusion.  \\subsection{Future Refinements}  There are several aspects of the coupled model that we wish to improve upon in the future. These include both dynamics and chemistry. Even though our spherical model captures the essence of the ambipolar diffusion-driven cloud evolution, the effects of magnetic tension (which tend to flatten a cloud) cannot be treated. Indeed, the opposite extreme, a disk-like geometry, may describe the observed elongated mass distribution in L1544 better. It should be straightforward to extend the model to this geometry, which would allow for the inclusion of rotation whose presence has been inferred for L1544 by Ohashi et al. (1999) based on CCS data. On the chemistry side, a major uncertainty is the adsorption energies for CO and other species. Different sets of adsorption energies have been considered by Aikawa et al. (2001), and a similar parameter study is needed for our model. Also uncertain are the probability of molecules sticking onto dust grains and grain surface processes, and an exploration of different cases is desirable. Another area of improvement would be the treatment of magnetic coupling coefficient, by including the effects of charged dust grains (especially small ones) and external UV radiation field."
    },
    "0201/astro-ph0201533_arXiv.txt": {
        "abstract": "The X-ray luminous RS CVn binary system HR~1099 has been observed on several occasions in the early phases of \\textit{Chandra} and \\textit{XMM-Newton}. A very hot (up to 40~MK) dominant coronal plasma has been identified from the high-resolution spectroscopic data; cooler plasma is seen down to about 3~MK. We recently obtained 100 ksec  in \\textit{Chandra}'s Guest Observer Program to study the corona of HR 1099 with the High-Resolution Camera (HRC-S) and the Low Energy Transmission Grating (LETG) across the complete temperature range above 1~MK.  The data provide an unprecedented  view of spectral lines and continua at high resolution between ~1 and 175~\\AA. We present our investigations on the \\textit{Chandra} LETGS observation of HR~1099 into the context of the latest results obtained with \\textit{XMM-Newton}. ",
        "introduction": "\\object{HR~1099} (=\\object{V711 Tau}; d=28.97~pc, K1~IV+G5~V-IV) is one of the brightest binary systems of the RS CVn class.  The active K subgiant mainly contributes to the chromospheric and coronal emissions (e.g., \\cite{maudard-B1-9:ayres01}). While in the solar corona, a First Ionization Potential (FIP) effect is observed (low-FIP elements enhanced by factors of 4-6 relative to high-FIP elements), an inverse FIP effect in  the corona of HR~1099 has been found (\\cite{maudard-B1-9:brinkman01}). Other active stars also show such an IFIP effect (\\cite{maudard-B1-9:guedel01a,maudard-B1-9:guedel01b,maudard-B1-9:drake01,maudard-B1-9:huenemoerder01}), while it is not present in the intermediately active Capella (\\cite{maudard-B1-9:audard01a}). It is unclear whether there is an enrichment or a depletion of the coronal material compared to the photospheric material in RS CVn systems, because of the high levels of activity and the short rotation periods. Estimates for HR~1099 range from [Fe/H]$=-0.6$ to 0. Hence the IFIP effect may only reflect the photospheric composition. However, in an analysis of solar analogs with known photospheric abundances (close to solar), \\cite*{maudard-B1-9:guedel02} (see also \\cite{maudard-B1-9:audard02} in these proceedings) find an evolution from an IFIP to a normal FIP effect with decreasing activity, suggesting that the transition is real.  \\begin{figure*} \\includegraphics[width=0.5\\linewidth]{maudard-B1-9_fig1a.eps} \\hfill \\includegraphics[width=0.5\\linewidth]{maudard-B1-9_fig1b.eps} \\vspace*{1mm} \\caption{LETGS spectrum of HR 1099 together with major identified emission lines (at the exception of the {\\rm Fe}~\\textsc{ix} and {\\rm Fe}~\\textsc{x} lines, in violet, that have been marked to emphasize their absence in the spectrum as an evidence of the absence of a significant cool component, $T \\approx 1 - 2$~MK, in the corona of HR~1099). Data are in red, and a best-fit model (cf. Fig.~\\ref{maudard-B1-9_fig:fig3}) is in green.} \\label{maudard-B1-9_fig:fig1} \\end{figure*} ",
        "conclusions": "HR~1099 is a hot RS CVn binary that shows a high Ne/Fe ratio. Relative to solar photospheric abundances, the abundance pattern of HR~1099 (normalized to O) is similar to other active RS CVn binaries, although less active stars ($\\lambda$~And, Capella) show no inverse FIP effect. This pattern in active binaries fits well into the long-term evolution from IFIP to FIP found in solar analogs (G\\\"udel et al. 2002; see also \\cite{maudard-B1-9:audard02}).  \\begin{figure}[!ht] \\includegraphics[width=\\linewidth]{maudard-B1-9_fig2.eps} \\caption{X-ray light curve of the LETGS observation of HR~1099. The sum of the positive and negative grating orders is shown. The light curve suggests that the observation was possibly performed during the decay phase of a flare.} \\label{maudard-B1-9_fig:fig2} \\end{figure}  \\begin{figure} \\centering \\includegraphics[width=\\linewidth]{maudard-B1-9_fig3.eps} \\caption{Emission measure distribution of HR~1099 observed by \\textit{Chandra} LETGS. Abundances from a multi-temperature fit have been used. Note that a modified MEKAL plasma emission code available in SPEX has been used.} \\label{maudard-B1-9_fig:fig3} \\end{figure}   \\begin{figure}[!ht] \\centering \\includegraphics[width=\\linewidth]{maudard-B1-9_fig4.eps} \\vspace*{5mm} \\caption{Coronal abundances of HR~1099 (normalized to oxygen) relative to solar photospheric abundances (Anders \\& Grevesse 1989, except {\\rm Fe}, Grevesse \\& Sauval 1999) as a function of FIP. Note that the abundances have been derived using a recent update of the MEKAL code in SPEX. Similarly to Brinkman et al. (2001), a high Ne/Fe abundance is found, suggesting the presence of an inverse FIP effect.} \\label{maudard-B1-9_fig:fig4} \\end{figure}  \\begin{figure} \\centering \\includegraphics[width=\\linewidth]{maudard-B1-9_fig5.eps} \\caption{Density upper limit from {\\rm Fe}~\\textsc{xxii} lines. The dotted blue line represents the measured line ratio, while the dashed blue lines are 1~$\\sigma$ confidence ranges. The continuous red curve is a theoretical curve based on models derived by Brickhouse et al. (1995). The black dash-dotted line gives the upper limit to the measured density. Note that the {\\rm O}~\\textsc{vii} line ratio (sensitive to plasma $\\approx 2$~MK) gives $n_\\mathrm{e} = 2 - 5 \\times 10^{10}$~cm$^{-3}$.} \\label{maudard-B1-9_fig:fig5} \\end{figure}  \\begin{figure*}[!h] \\centering \\includegraphics[width=0.75\\linewidth]{maudard-B1-9_fig6.eps} \\caption{\\textit{XMM-Newton} RGS spectra of HR~1099 with compared other RS CVn binaries (see Audard \\& G\\\"udel 2002 in these proceedings).  Note the difference in flux ratios between {\\rm Ne}~\\textsc{ix} (log$T_{\\mathrm{m}} \\approx 6.6$; arrow) and {\\rm Fe}~\\textsc{xvii} (log$T_{\\mathrm{m}} \\approx 6.73$; arrow), especially in UX Ari and Capella, suggesting that their coronae differ in their elemental composition.} \\label{maudard-B1-9_fig:fig6} \\end{figure*}"
    },
    "0201/astro-ph0201475_arXiv.txt": {
        "abstract": "I review the observational evidence for stream-fed accretion in intermediate polars. Recent work on the discless system V2400~Oph confirms the pole-flipping model of stream-fed accretion, but this applies only to a minority of the flow. The bulk of the flow is in the form of blobs circling the white dwarf, a state which might have been a precursor to disc formation in other IPs. I also discuss work on the systems with anomalously long spin periods, V1025~Cen and EX~Hya. There are arguments both for and against stream-fed accretion in V1025~Cen, and further work is necessary before reaching a conclusion about this system. ",
        "introduction": "Early work on intermediate polars (IPs), taking a lead from studies of DQ~Her, tended to assume the presence of an accretion disc. The accretion stream would feed the outer disc, as in non-magnetic systems, and the magnetic field would disrupt the inner disc and so channel the accretion along field lines onto magnetic polecaps. Hameury, King, \\&\\ Lasota (1986) and Lasota \\&\\ King (1991) presented theoretical arguments against this idea. They suggested that most IPs were discless, with the accretion stream falling until it encountered the magnetosphere directly.  They argued that, observationally, discless systems would be distinguished by X-ray modulations over the orbital cycle, since in a stream-fed system the accretion sites would lie `beneath' the stream, and so be localised in orbital phase.  Hellier (1991) and Hellier, Garlick, \\&\\ Mason (1993) pointed out that X-ray orbital modulations can also result from obscuration of the white dwarf by structure in an accretion disc, as seen in `dipping' LMXBs, and so are not exclusive to discless systems.  We suggested that a better diagnostic was the presence of an X-ray pulsation at the beat frequency ($\\omega$\\,--\\,$\\Omega$) between the spin ($\\omega$) and orbital ($\\Omega$) frequencies.  This arises because the geometry changes on the beat frequency when a magnetosphere (spinning at $\\omega$) rotates beneath a stream (orbiting at $\\Omega$). Indeed we would expect the stream to follow the field line involving least deviation from the orbital plane, and so flow to the upper magnetic pole for half the beat cycle and then flip to the lower pole for the remainder of the cycle.  Wynn \\&\\ King (1992) modelled the situation and confirmed that X-ray modulations at $\\Omega$ and $\\omega$\\,--\\,$\\Omega$ are characteristic of stream-fed accretion, but also found that power could be moved from $\\omega$\\,--\\,$\\Omega$ to 2$\\omega$\\,--\\,$\\Omega$ for certain combinations of the system inclination and the angle between the magnetic and spin axes.  Hellier (1991; 1992) analysed X-ray lightcurves to show that all currently known IPs were dominated by X-ray pulsations at the spin period, rather than the beat period, and so argued that they were all disc-fed accretors. However, Buckley et\\,al.\\ (1995; 1997) then discovered the {\\it Rosat\\/} source V2400~Oph, which showed a strong X-ray pulsation at the 1003-sec beat period but none at the 927-sec spin period, and thus was the first secure case of discless accretion amongst the known IPs.  The 927-sec period was detected only in polarised light, and this raised the following problem for deductions based on X-ray periodicities: since many IPs show no polarised light, might we fail to detect the spin period in a non-polarised discless system, and if so might we be misinterpreting the beat cycles as spin cycles? Exactly this issue was debated in the case of BG~CMi (Norton et\\,al.\\ 1992b; de Martino et\\,al.\\ 1995; Hellier 1997).  New spectroscopy of V2400~Oph, reported in Hellier \\&\\ Beardmore (2002), is reassurring on this point.  Both the spin and the beat cycles are seen easily in the emission lines (Fig.~1), and this would leave little doubt as to their respective identities, even if we had no polarimetry.  Thus V2400~Oph greatly strengthens the overall argument, since the discovery of a system showing so clearly the predicted hallmark of discless accretion confirms the validity of distinguishing between disc-fed and stream-fed accretors using the spin and beat pulses in the X-ray lightcurve.  \\begin{figure}[t]   % \\vspace*{8cm} \\special{psfile=hgottf1.eps hscale=67 vscale=67 hoffset=0 voffset=0 angle=0} \\caption{Fourier transforms from V2400~Oph, comparing the polarised light (data by David Buckley) with X-ray data and H$\\beta$ V/R ratios. The spin, orbital and beat frequencies are marked with the usual notation of $\\omega$, $\\Omega$ and $\\omega$\\,--\\,$\\Omega$, respectively. Peaks marked `sp' are windowing caused by the spacecraft orbit.} \\end{figure}  Since nearly all well-studied IPs show both spin and beat cycles in the emission lines, we can have confidence in assigning cycles correctly and thus in applying the argument to the whole class. ",
        "conclusions": ""
    },
    "0201/astro-ph0201369_arXiv.txt": {
        "abstract": " ",
        "introduction": "One of the most exciting aspects of modern astrophysics is the possible existence of a new family of compact stars, which are made entirely of deconfined {\\it u,d,s} quark matter ({\\it strange quark  matter} (SQM)). These strange quark matter stars are called in the scientific literature {\\it strange stars} (SS). They differ from neutron stars, where quarks are confined within neutrons, protons, and eventually within other hadrons (hadronic matter stars). The possible existence of SS is a direct consequence of the so called {\\it strange matter hypothesis} \\cite{witt}. According to this hypothesis, SQM (in equilibrium with respect to the weak interactions) could be the true ground state of matter. In other words, one assumes the energy per baryon of SQM (at the baryon density where the pressure is equal to zero) to be less than the lowest energy per baryon found in nuclei, which is about 930 {\\rm MeV} for $^{56}$Fe. According to the strange matter hypothesis, the ordinary state of matter, in which quarks are confined within hadrons, is a metastable state \\cite{witt,fj84}. The strange matter hypothesis does not conflict with the existence of atomic nuclei as conglomerates of nucleons, or with the stability of ordinary matter \\cite{fj84,mads99,bomb2001}.  From a basic point of view the equation of state for SQM should be calculated solving QCD at finite density. As we know, such a fundamental approach is presently not doable. Therefore one has to rely on phenomenological models. In this work, we use two simple phenomenological models for the equation of state (EOS) of strange quark matter. One is a model \\cite{fj84} which is related to the MIT bag model for hadrons. The other is a model proposed by Dey {\\it et al.} \\cite{dey98}. ",
        "conclusions": ""
    },
    "0201/astro-ph0201182_arXiv.txt": {
        "abstract": "We present a series of simulations to demonstrate that high-fidelity velocity-delay maps of the emission-line regions in active galactic nuclei can be obtained from time-resolved spectrophotometric data sets like those that will arise from the proposed \\Kronos\\ satellite. While previous reverberation-mapping experiments have established the size scale $R$ of the broad emission-line regions from the mean time delay $\\tau = R/c$ between the line and continuum variations and have provided strong evidence for supermassive black holes, the detailed structure and kinematics of the broad-line region remain ambiguous and poorly constrained. Here we outline the technical improvements that will be required to successfully map broad-line regions by reverberation techniques. For typical AGN continuum light curves, characterized by power-law power spectra $P(f) \\propto f^{-\\alpha}$ with $\\alpha = -1.5\\pm0.5$, our simulations show that a small UV/optical spectrometer like \\Kronos\\ will clearly distinguish between currently viable alternative kinematic models. From spectra sampling at time intervals \\Tres\\ and sustained for a total duration \\Tdur, we can reconstruct high-fidelity velocity-delay maps with velocity resolution comparable to that of the spectra, and delay resolution $\\Delta \\tau \\approx 2 \\Tres$, provided $\\Tdur$ exceeds the BLR crossing time by at least a factor of three. Even very complicated kinematical models, such as a Keplerian flow with superimposed spiral wave pattern, are cleanly resolved in maps from our simulated \\Kronos\\ datasets. Reverberation mapping with \\Kronos\\ data is therefore likely deliver the first clear maps of the geometry and kinematics in the broad emission-line regions 1--100 microarcseconds from supermassive black holes. ",
        "introduction": "Reverberation mapping (Blandford \\& McKee 1982) has become part of the standard toolkit for investigation of active galactic nuclei (AGNs). Detailed comparison of broad emission-line flux variations and the continuum variations that drive them can be used to determine the structure and kinematics of the broad-line region (BLR) under a rather simple and straightforward set of assumptions (see Horne 1999 and Peterson 2001 for primers on reverberation mapping). AGN emission lines are observed to respond roughly linearly to continuum variations, each line having a different range of time delays, $\\tau$. Since reprocessing times are relatively short, the time delays are dominated by light travel time, \\begin{equation} \\label{eq:delay} \\tau = \\frac{R}{c}( 1 + \\cos{\\theta} ) \\ , \\end{equation} where ($R$,$\\theta$,$\\phi$) are spherical polar coordinates with $\\theta=0$ along our line-of-sight beyond the compact continuum source at $R=0$. Iso-delay surfaces slice up the geometry of the emission-line region on a set of nested paraboloids. In the simplest linear reprocessing model, the relationship between the continuum light curve, $F_c(t)$, and the emission-line profile variations, $F_\\ell(v,t)$, is given by \\begin{equation} \\label{eq:linear} F_\\ell(v,t) = \\int \\Psi_\\ell(v,\\tau) F_c(t - \\tau)\\, d\\tau \\ . \\end{equation} The ``transfer function,'' or ``velocity-delay map'', \\begin{equation} \\label{eq:psi} \\Psi_\\ell(v,\\tau) = \\frac{1}{d\\tau} \\frac{ \\displaystyle \\partial F_\\ell(v,t) }{ \\displaystyle \\partial F_c(t - \\tau) } \\ , \\end{equation} is the response at velocity $v$ and time delay $\\tau$ in the flux of emission line $\\ell$. The velocity-delay maps of various emission lines code information on the geometry, kinematics, and physical conditions in the BLR. The immediate goal of reverberation mapping is to recover $\\Psi_\\ell(v,\\tau)$ for each emission line by detailed fitting to variability recorded in high-quality time-resolved spectrophotometric data. The ultimate goal is to unravel the geometry, kinematics, and physical conditions in the BLR.  To date, the great success of reverberation mapping has been determination of the mean broad-line response times for approximately three dozen AGNs, in several cases for multiple lines in a single source (see the compilation of Kaspi et al.\\ 2000). By combining the response times (or ``lags'') with line widths and assuming that gravity is the dominant force acting on the line-emitting gas, a virial mass for the central black hole can be deduced. Two lines of evidence suggest that these reverberation-based masses are reasonably accurate and actually do measure the black hole mass: \\begin{enumerate} \\item For AGNs in which the lags and widths of multiple lines have been measured, there is a clear anticorrelation between Doppler line width $\\vres$ and response time $\\tau$ that is consistent with the virial prediction $\\tau \\propto \\vres^{-2}$ since $\\tau$ measures the light-travel time across the BLR (Peterson \\& Wandel 1999, 2000; Onken \\& Peterson 2002). Lines that arise in higher ionization-level gas (e.g., \\heii\\,$\\lambda1640$, \\nv\\,$\\lambda1240$) are broader and have shorter response times than lines that arise primarily in lower ionization-level gas (e.g., \\Hbeta, \\ciii]\\,$\\lambda1909$). \\item Comparison of reverberation-based black-hole masses and host-galaxy bulge velocity dispersons shows a relationship that appears to be the same as the black-hole/bulge velocity-dispersion relationship seen in quiescent galaxies (Ferrarese et al.\\ 2001). \\end{enumerate}  On the other hand, reverberation mapping has not yet achieved the motivating design goal of determining the actual geometry and velocity field of the BLR. This is not surprising since inversion of eq.\\ (\\ref{eq:linear}) to solve for the velocity-delay map requires large amounts of high-quality spectra, which are difficult to obtain for AGNs. Published attempts include analyses of NGC~4151 (Ulrich \\& Horne 1996), and NGC~5548 (Wanders et al.\\ 1995; Done \\& Krolik 1996), both of which yielded rather ambiguous results. Nevertheless, pursuit of the science goal of determining the BLR structure is important: \\begin{enumerate} \\item The origin of the BLR emission remains one of the major mysteries of AGNs. There is no widely accepted paradigm for the nature of the BLR. The BLR clouds may represent the cooler and denser component of a two-phase medium in pressure and virial equilibrium (Krolik, McKee \\& Tartar 1981). Alternatively, the BLR clouds might be in virial motion, but magnetically confined (Rees 1987). Compelling arguments have been made that the broad emission arises as part of a massive outflow (e.g., Chiang \\& Murray 1996; Bottorff et al.\\ 1997). Others have suggested that the broad emission arises in the extended atmospheres of stars (e.g., Alexander \\& Netzer 1997) or at least in part from the surface of the accretion disk itself (e.g., Collin-Souffrin et al.\\ 1988). \\item Knowledge of the BLR structure and kinematics is necessary to understand the potential systematic uncertainties in reverberation-based black-hole masses. This is especially interesting if indeed the outflow models are correct since the virial relationship $\\tau \\propto \\vres^{-2}$ is unexpected in simple outflow models. Reverberation-based black-hole masses are now being used to anchor secondary methods of mass determination (e.g., Laor 1998; Wandel, Peterson, \\& Malkan 1999; Vestergaard 2002) and it is therefore essential to understand how the measurements are affected by systematics. \\end{enumerate}  Successful reverberation mapping requires a combination of high time resolution, long duration, homogenous high signal-to-noise data, and reasonably high spectral resolution. Precise specification of these quantities depends on physical time scales of the source (e.g., BLR light-crossing time) and the character of the variations, which are neither simple nor regular. This makes experimental design difficult. Nevertheless, the first generation of reverberation experiments have provided a great deal of insight on the continuum and emission-line variability properties of AGNs, making it possible to model their behavior through numerical simulations. Detailed simulations can then be used to determine the observational requirements to map the BLR.  Our main goal in this paper is to describe a program we have undertaken to define the requirements for recovery of high-fidelity velocity-delay maps from data obtained with a small UV/optical spectrophotometer. For the sake of realism, we have used as a model detector system that of a proposed multiwavelength observatory, \\Kronos\\ (Polidan \\& Peterson 2001). The wavelength range and resolution and achievable signal-to-noise ratios ($S/N$) in these simulations are set by the specifications for \\Kronos.  We have carried out two series of simulations, a comparatively simple first series, and a more realistic and complicated second series. The first series of simulations, discussed in \\S\\ref{sec:series1}, addresses the use of velocity-delay maps for a single line, for example the \\civ\\ line, to distinguish among several alternative and currently viable models for the geometry and kinematics of the BLR.  The second series, discussed in \\S\\ref{sec:series2}, is intended to be more realistic and challenging. We chose a model with a complicated yet plausible structure, specifically an inclined Keplerian disk with a 2-armed spiral density wave superposed on it. We then used a photoionization equilibrium model appropriate for NGC~5548 (Kaspi \\& Netzer 1999) to compute the anisotropic emissivity in each line at each point in the disk. In this series of simulations, we considered the response of multiple lines and constructed a model spectrum at each time to determine how blending of closely spaced lines (e.g., \\Lya\\,$\\lambda1216$ and \\nv\\,$\\lambda1240$) affect the experimental results. ",
        "conclusions": "The major result of this investigation is  a clear demonstration that with technically realizable observational programs reverberation mapping can successfully recover even complex emission-line velocity-delay maps.  Previous reverberation mapping programs have had comparatively modest goals: generally, the intent has been to measure the mean response time for various emission lines. From these programs, we have learned enough about AGN continuum and emission-line variability characteristics to carry out realistic simulations, such as those described here, that will define future programs with more ambitious goals. These simulations show clearly that even the most ambitious previous programs could not be expected to yield the results we now seek, i.e., a complete velocity-delay map that can be used to identify the detailed structure and kinematics of the BLR.  We also conclude from these simulations that while high-fidelity reverberation mapping of even a single line will be a tremendous step forward, this will not yield the complete structure of the BLR. This is clearly illustrated in Fig.\\ \\ref{fig:4}, where the vivid colors arise because different emission lines probe different ranges of physical conditions and ionization level. To acquire a complete picture of the BLR, a variety of lines spanning a broad range of ionization level (and hence mean response time) need to be mapped. Moreover, it is distinctly possible that the high-ionization and low-ionization lines arise primarily in physically distinct regions with different geometries and kinematics (e.g., Collin-Souffrin et al.\\ 1988); if this is true, the need for reverberation mapping of multiple emission lines is self-evident.  Once we have acquired the data to make high-fidelity velocity-delay maps for different emission lines, how can we produce a map of the BLR? Perhaps the simplest approach is though inspired modeling: practitioners recognizing the structure in the velocity-delay maps can devise appropriate models with adjustable parameters to obtain the best fit of a model to the data. A more ambitious program aims to reconstruct a complete phase-space map from the velocity-delay maps of different lines, or even directly from the data. This may be problematic, however, because the ``observable'' velocity-delay map is a two-dimensional projection of the six-dimensional phase space, so there are degeneracies. These can be partially resolved by combining results from multiple lines with photoionization equilibrium models, and some success has already been achieved with simple geometries (Horne 2001, Horne, Korista, \\& Goad 2002). Such methods will succeed best if the BLR structure has some degree of simplifying symmetry, as we expect from most current models. But even if the BLR is completely chaotic with no symmetries, this could be concluded from good reverberation data, and would provide us with an important answer about the inner structure of AGNs."
    },
    "0201/astro-ph0201461_arXiv.txt": {
        "abstract": "On September 21 at 18950.56 SOD (05:15:50.56) UT the FREGATE $\\gamma$-ray instrument on the High Energy Transient Explorer (HETE) detected a bright gamma-ray burst (GRB).  The burst was also seen by the X-detector on the WXM X-ray instrument and was therefore well-localized in the X direction; however, the burst was outside the fully-coded field-of-view of the WXM Y-detector, and therefore information on the Y direction of the burst was limited. Cross-correlation of the HETE and {\\it Ulysses} time histories yielded an Interplanetary Network (IPN) annulus that crosses the HETE error strip at a $\\sim$45 degree angle.  The intersection of the HETE error strip and the IPN annulus produces a diamond-shaped error region for the location of the burst having an area of 310 square arcminutes.  Based on the FREGATE and WXM light curves, the duration of the burst is characterized by a $t_{90}$ = 18.4 s in the WXM 4 - 25 keV energy range, and 23.8 s and 21.8 s in the FREGATE 6 - 40 and 32 - 400 keV energy ranges, respectively.  The fluence of the burst in these same energy ranges is 4.8 $10^{-6}$, 5.5 $10^{-6}$, and 11.4 $10^{-6}$ erg cm$^{-2}$, respectively. Subsequent optical and radio observations by ground-based observers have identified the afterglow of GRB010921 and determined an apparent redshift of z = 0.450. ",
        "introduction": "\\setcounter{footnote}{0}  As has long been recognized, accurate locations and rapid follow-up observations in many wavelengths are central to understanding the nature of gamma-ray bursts. For this reason, a strategy evolved in the late 1970's and early 1980's (e.g., Woosley et al. 1982) to detect GRBs, not only in gamma-rays, but also in emitted X-rays and optical light. While detection of short transients at the lower energies posed observational challenges, and the strength of the optical signal was unknown, the possibility of arc minute localizations deduced from the X-rays that were known to be present was very appealing. This strategy was implemented in the HETE-1 (High Energy Transient Explorer) satellite, which was unfortunately lost due to a rocket failure on 1996 November 4, and in the highly successful \\textit{BeppoSAX} Mission (Costa et al. 1997).  The HETE-2 satellite (henceforth simply ``HETE''), which was successfully launched into equatorial orbit on 9 October 2000, is the first space mission entirely devoted to the study of gamma-ray bursts (GRBs). HETE utilizes a matched suite of low energy X-ray, medium energy X-ray, and gamma-ray detectors mounted on a compact spacecraft. A unique feature of HETE is its capability for localizing GRBs with $\\sim$1-10{\\arcmin} accuracy in real time aboard the spacecraft. GRB locations are transmitted, within seconds to minutes, directly to a dedicated network of telemetry receivers at 13 automated ``Burst Alert Stations\" (BAS) sited along the satellite ground track (Villasenor et al 2002). The BAS network then re-distributes the GRB locations world-wide to all interested observers via Internet and the GRB Coordinates Network (GCN) in $\\approx$1 s (Vanderspek et al. 2002, Barthelmy et al. 2002).  Thus, prompt optical, IR, and radio follow-up identifications can be anticipated for a large fraction of HETE GRBs.  Here we report the localization of the first HETE-discovered \\footnote{HETE detects $\\sim$50 GRBs yr$^{-1}$, of which $\\sim$15 yr$^{-1}$ are localized by the WXM (Ricker et al. 2002)} GRB for which a counterpart has been found. Based on the combined data from HETE and the Interplanetary Network (IPN), a diamond-shaped error box roughly 15 arc minutes on an edge was established for GRB010921. Subsequent searches of this error box by ground-based optical and radio instruments revealed a fading counterpart whose properties are reported elsewhere (Price et al 2001a, 2002; Kulkarni et al 2002; Park et al 2001a, 2002; Djorgovski et al. 2001). ",
        "conclusions": "The localization of GRB010921 by HETE and the IPN  has led directly to a successful identification of a low-redshift afterglow counterpart. The prompt initial GCN alert for GRB010921 was distributed worldwide within 17 s of the burst; however, the determination of the HETE WXM localization was delayed by 5.1 hours due to the need to prototype new software to accommodate GRB010921's extreme off-axis position in one of the WXM cameras. A confirmation, and definitive refinement of the WXM localization, using IPN data available 15.2 hours after the GRB enabled ground-based observers to successfully target GRB010921 during the first night following the burst, and to establish a counterpart well within the HETE-IPN error region. Hopefully, over the coming months the identification of optical counterparts resulting from prompt ($\\le$100 s), accurate HETE localizations will become routine. With the currently-projected long orbital lifetime ($>$10 years) and excellent health of the HETE spacecraft and instruments, we look forward to providing a uniquely valuable service to the worldwide community of GRB observers."
    },
    "0201/astro-ph0201527_arXiv.txt": {
        "abstract": "We have run Monte Carlo simulations, for quasar clustering redshift distortions in the Two-Degree Field QSO Redshift Survey (2QZ), in order to elicit the power of redshift distortions (geometric Alcock-Paczy\\'nski and linear kinematic) to constrain the cosmological density and equation of state parameters, $\\Omega_{m0}, \\Omega_{x0}, w$, of a pressureless matter + dark energy model. It turns out that, for the cosmological constant case ($w=-1$), the test is especially sensitive to the difference $\\Delta:=\\Omega_{m0}-\\Omega_{\\Lambda 0}$, whereas for the spatially flat case ($k=0$), it is quite competitive with SNAP and DEEP, besides being complimentary to them; furthermore, we find that, whereas not knowing the actual value of the bias does not compromise the correct recovering of $\\Delta$, taking into account the linear velocity effect is absolutely relevant, all within the $2\\sigma$ confidence level. ",
        "introduction": "The new millennium has ushered in a golden era for cosmology, driven by a flood of high quality observational data, from supernovae \\cite{Riess99,Perlmutter99,snap} to cosmic microwave background \\cite{deBernardis00,Balbi00,Pryke01,map,planck}, passing through galaxies and quasars \\cite{2df,sdss}, to mention just a few. All of them favor a spatially flat cosmological model, with a nonrelativistic matter with density parameter $\\Omega _{m0}\\simeq 1/3$ and a negative-pressure dark energy component with density parameter $\\Omega _{x0}\\simeq 2/3$. The exact nature, however, of this dark energy is not currently well understood, possible alternatives being a vacuum energy or cosmological constant ($\\Lambda$) or a dynamical scalar field (quintessence) \\cite{Ratra88,Frieman95,Caldwell98,Ferreira98}. An important task for present cosmology is thus to find new methods that can probe the amount of dark energy present in the Universe as well as its equation of state. These new methods may constrain distinct regions of the parameter space and are usually subject to different systematic errors.  The test we focus on here is the one suggested by Alcock and Paczy\\'nski (hereafter AP)\\cite{Alcock79}, which has attracted a lot of attention during the last years \\cite{Ryden95,Ballinger96,Matsubara96,Hui99,McDonald99,Kujat01}. This test is based on the fact that transverse (angular separation) and radial (redshift separation) distances have a different dependence on cosmological parameters, rendering a high redshift spherical object in real space distorted in redshift space. The degree of distortion increases with redshift and is very sensitive to $\\Lambda$ or, more generally, to dark energy. In particular, Popowski \\emph{et al.} \\cite{Popowski98} (hereafter PWRO) extended a calculation by Phillips \\cite{Phillipps94} of the geometrical distortion of the QSO correlation function. They suggested a simple Monte Carlo experiment to see what constraints should be expected from the 2dF QSO Redshift Survey (2QZ) and the Sloan Digital Sky Survey (SDSS). However, they did not estimate the probability density in the parameter space and, as a consequence, they could not notice that the test is in fact very sensitive to the difference $\\Omega_{m0}-\\Omega_{\\Lambda 0}$. Further, they did not take into account the effect of peculiar velocities, although they discussed its role arguing that it would not overwhelm the geometric signal.  Here we summarize the results which confirm the feasibility of redshift distortion (geometric AP + peculiar velocity) measurements to constrain cosmological parameters, by extending the PWRO Monte Carlo experiments and obtaining confidence regions in the ($\\Omega _{m0},\\Omega _{\\Lambda 0} $) and ($\\Omega_{m0}, w$) planes. We compare the expected constraints from the AP test, when applied to the 2QZ survey, with those obtained by other methods. We include a general dark energy component with equation of state $P_{x}=w\\,\\rho_{x}$, with $w$ constant. Our analysis can be generalized to dynamical scalar field cosmologies as well as to any model with redshift dependent equation of state. Since most quasars have redshift at around $z= 2$ we expect the test to be useful in the determination of a possible redshift dependence of the equation of state. We explicitly take into account the effect of large-scale coherent peculiar velocities. Our calculations are based on the measured 2QZ distribution function and we consider best fit values for the amplitude and exponent of the correlation function as obtained by Croom {\\it et al.} \\cite{Croom01}. In this work, we only consider the 2QZ survey although the results can easily be generalized to SDSS. ",
        "conclusions": "In Figure \\ref{fig:w-1}, we show the predicted AP likelihood contours in the ($\\Omega_{m0},\\Omega_{\\Lambda 0}$)-plane for the 2QZ survey (solid lines), in the case $w=-1$, in a universe with arbitrary spatial curvature. The scattered points represent maximum likelihood best fit values for $\\Omega_{m0}$ and $\\Omega_{\\Lambda 0}$. The assumed ``true'' values are ($\\Omega_{m0}=0.3,\\Omega_{\\Lambda 0}=0$) and ($0.28,0.72$), for the top and bottom panels, respectively. In the top panel the displayed curve corresponds to the predicted $2\\sigma$ likelihood contour. In the bottom panel the predicted $1\\sigma$ contour (dashed line) for one year of SNAP data \\cite{Goliath01} is displayed, together with the predicted $1\\sigma$ AP contour. For the SNAP contour, it is assumed that the intercept $\\cal{M}$ is exactly known. To have some ground of comparison with current SNe Ia observations, in the same panel, we also plot (dotted lines) the Supernova Cosmology Project \\cite{Perlmutter99} $1\\sigma$ contour (fit C). As expected, in both cases, the test recovers nicely the ``true'' values. We stress out that the test is very sensitive to the difference $\\Omega_{m0}-\\Omega_{\\Lambda 0}$. From the bottom panel we note that the sensitivity to this difference is comparable to that expected from SNAP, of the order $\\pm 0.01$. Comparatively, however, the test has a larger uncertainty in the determination of $\\Omega_{m0}+\\Omega_{\\Lambda 0}$, of the order $\\pm 0.17$. The degeneracy in $\\Omega_{m0}+\\Omega_{\\Lambda 0}$ may be broken if we combine the estimated results for the AP test with, for instance, those from CMB anisotropy measurements, whose contour lines are orthogonal to those exhibited in the panels \\cite{Hu99}. \\begin{figure} \\centering \\includegraphics[height=14cm]{fig1.eps} \\caption{Simulated models at fixed $w=-1$ and corresponding predicted AP confidence contours (solid lines). In the top panel we show the predicted $2\\sigma$ likelihood contour assuming a ``true'' model ($\\Omega_{m0}=0.3, \\Omega_{\\Lambda 0}=0)$. In the bottom panel the predicted $1\\sigma$ contour (dashed line) for one year of SNAP data \\protect\\cite{Goliath01} is displayed, together with the predicted $1\\sigma$ AP contour. For both tests we consider $\\Omega_{m0}=0.28$ and $\\Omega_{\\Lambda 0}=0.72$; also displayed is a $1\\sigma$ confidence contour obtained by the Supernova Cosmology Project (dotted lines; \\protect\\cite{Perlmutter99}).}\\label{fig:w-1} \\end{figure}  In order to estimate the consequences of neglecting the effect of linear peculiar velocities, in the top panel of Figure \\ref{fig:pvfavb}, we included them in the calculation of the $A_{i}$ values but neglected them in the computation of the maximum likelihood; in this panel, we assume $\\Omega _{m0}=0.3$ and $\\Omega _{\\Lambda 0}=0$ as ``true'' values. Notice that the point with the ``true'' $\\Omega _{m0}$ and $\\Omega _{\\Lambda 0}$ values is outside the $2\\sigma $ contour. It is clear, therefore, the necessity of taking this effect in consideration when analyzing real data.  To illustrate that the AP test is in fact more sensitive to the mean amplitude of the bias rather than to its exact redshift dependence, we plot, in the bottom panel of Figure \\ref{fig:pvfavb}, the $2\\sigma$ contour line, assuming as ``true'' values $\\Omega_{m0}=0.3$ and $\\Omega_{\\Lambda 0}=0.7$. For this panel, the ``true'' $A_i$ values were generated assuming $b_0=1.45$ and $m=1.68$. However, for the simulations, we considered a constant bias ($m=0$), such that $b_{0,sim}:=\\int_{z=z_{min}}^{z_{max}}F(z)b_{true}(z)dz=2.46$. We remark that the contour is slightly enlarged, mainly in the direction of the ``ellipsis'' major axis. However, the uncertainty in $\\Omega_{m0}-\\Omega_{\\Lambda 0}$ is practically unaltered, confirming the strength of the test \\cite{Yamamoto01}. We did the same analysis assuming $\\Omega_{m0}=1$ and $\\Omega_{\\Lambda 0}=0$ and obtained similar results.  \\begin{figure} \\centering \\includegraphics[height= 14cm]{fig2.eps} \\caption{Simulated models at fixed $w=-1$ and corresponding $2\\sigma$ predicted AP confidence contour; in both panels, the ``true'' model is indicated by a solid dot. Top panel: The ``true'' model, $(0.3, 0)$, takes into account the effect of peculiar velocities, but the simulated ones do not. Notice that the ``true'' model does not fall into the $2\\sigma$ confidence region. Bottom panel: The ``true'' model, $(0.3, 0.7)$, uses a redshift dependent bias function with $b_0=1.45$ and $m=1.68$, whereas the simulated ones use a constant bias equal to 2.46.}\\label{fig:pvfavb} \\end{figure}  In Figure \\ref{fig:k0}, we show the predicted AP likelihood contours in the $(\\Omega_{m0},w)$-plane for the 2QZ survey (solid lines) for flat models ($\\Omega_{k0}=0$). The ``true'' values are $(\\Omega_{m0}=0.28,w=-1)$ and $(\\Omega_{m0}=0.3,w=-0.7)$ for the top and bottom panels, respectively. In the top panel, we show, besides the AP contour, the predicted contour for one year of SNAP data (dashed line; \\cite{Goliath01}), both at $1\\sigma$ level. For the SNAP contour, the intercept $\\cal{M}$ is assumed to be exactly known. Notice that the contours are somewhat complementary and are similar in strength. In the bottom panel, we compare the predicted $95\\%$ confidence contour of the AP test with the same confidence contour for the number count test as expected from the DEEP redshift survey (dashed line; \\cite{Newman00}). Again the contours are complementary, but the uncertainties on $\\Omega_{m0}$ and $w$ for the AP test are quite smaller.  \\begin{figure}[htb] \\centering \\includegraphics[height=14cm]{fig3.eps} \\caption{Simulated flat models and corresponding predicted AP confidence contours (solid lines). The top panel is from a ``true'' model ($\\Omega_{m0}=0.28$, $w=-1$), and displays the predicted confidence contours for the AP test and the SNAP mission (dashed line; \\protect\\cite{Goliath01}), both at $1\\sigma$ level. The bottom panel is from a ``true'' model ($\\Omega_{m0}=0.3$, $w=-0.7$), and displays the predicted confidence contours for the AP test and the DEEP survey (dashed line; \\protect\\cite{Newman00}), both at the $95\\%$ level.}\\label{fig:k0} \\end{figure}  In summary, we have shown that the Alcock-Paczy\\'nski test applied to the 2dF quasar survey (2QZ) is a potent tool for measuring cosmological parameters. We stress out that the test is especially sensitive to $\\Omega_{m0}-\\Omega_{\\Lambda 0}$. We have established that the expected confidence contours are in general complementary to those obtained by other methods and we again emphasize the importance of combining them to constrain even more the parameter space. We have also revealed that, for flat models, the estimated constraints are similar in strength to those from SNAP with the advantage that the 2QZ survey will soon be completed.  Of course our analysis can be improved in several aspects. For instance, for the fiducial Einstein-de Sitter model, we have assumed that $\\gamma$ and $r_0$ do not depend on redshift. In fact, observations \\cite{Croom01} seem to support these assumptions, but further investigations are necessary. Since the test is very sensitive to $\\Omega_{m0}-\\Omega_{\\Lambda 0}$, the effect of small-scale peculiar velocities should also be incorporated in future analyses in order to eliminate any potential source of systematic bias. At present, the quasar clustering bias is not completely well understood. Theoretical as well as observational progress in its determination will certainly improve the real capacity of the test. However, confirming previous investigations \\cite{Yamamoto01}, we have found that the test is, in fact, more sensitive to the mean amplitude of the bias rather than to its exact redshift dependence. A more extensive detailed report of this work can be found in \\cite{Calvao02}."
    },
    "0201/hep-ph0201154_arXiv.txt": {
        "abstract": "We study the emissivity properties of a geometrically thin, optically thick, steady accretion disc about a static boson star. Starting from a numerical computation of the metric potentials and the rotational velocities of the particles in the vicinity of the compact object, we obtain the power per unit area, the temperature of the disc, and the spectrum of the emitted radiation. In order to see if different central objects could be actually distinguished, all these results are compared with the case of a central Schwarzschild black hole of equal mass. We considered different situations both for the boson star, assumed with and without self-interactions, and the disc, whose internal commencement can be closer to the center than in the black hole case. We finally make some considerations about the Eddington luminosity, which becomes radially dependent for a transparent object. We found that, particularly at high energies, differences in the emitted spectrum are notorious. Reasons for that are discussed.\\\\  PACS Number(s): 04.40.Dg, 98.62.Mw, 04.70.-s ",
        "introduction": "\\indent The possible existence of very massive non-baryonic objects in the center of some galaxies is being studied since a long time. As far as we are now aware, they were first hypothesized by Tkachev \\cite{TKA}, who also studied the emissivity properties arising from particle anti-particle annihilation processes. More recently, detailed studies of the properties of neutrino ball scenarios have also been carried out by Viollier and his collaborators \\cite{NB}. In Ref. \\cite{GALAXY}, in addition, we have also explored whether supermassive non-baryonic boson stars might be the central object of some galaxies. To fix the situation to a particular case, we have paid special attention to the Milky Way. This study had a twofold aim. On one hand, it focused on what current dynamical observational data have established regarding the properties of the galactic center. We have concluded in this sense that scalar stars fitted very well into these dynamical constraints. On the other hand, we have also discussed what kind of observations could actually distinguish between a supermassive black hole and a boson star of equal mass, pointing out several possible tests.\\\\  In the case of our own Galaxy, recent observations probe the gravitational potential at a radius larger than $\\sim 10^{4}$ Schwarzschild radii \\cite{Genzel}. The black hole scenario is the current paradigm, but suggestions that the dark central objects are indeed black holes are based only on indirect astrophysical arguments, basically dynamical in nature, that could be sustained by any small relativistic object other than a black hole, if it exists \\cite{kor95}. It is then advisable to explore possible alternative scenarios. The final aim should be to devise definitive tests that could observationally solve the issue. One is to look for the event horizon projected onto the sky plane. Although this is the key concept of both, Constellation-X \\cite{cons} and Maxim \\cite{maxim} satellites, the needed spatial resolution is still a dream for the future. It might be more feasible to look for the realization of the recently introduced concept of the shadow of a black hole \\cite{shadow}. Although any highly relativistic object would also produce a shadow, the observed features might be enough to decide whether an event horizon is present or not. \\\\  From the point of view of boson stars physics (see \\cite{MIELKE} for a recent review), we also need to know whether fundamental scalars capable to form the stars do exist; observational tests in this sense are highly desirable too. Only after the discovery of the boson mass spectrum we shall be in position to determine which galaxies, if any at all, could be modeled with such a center. In recent works, Schunck and Liddle \\cite{sch97}, Schunck and Torres \\cite{sch00}, and Capozziello et al. \\cite{CAP00}, among others, analyzed different observational effects that boson stars would produce. Also, boson stars were proposed as sources for some of the gamma ray bursts \\cite{IWA}, and as a possible lens in a gravitational lensing configuration \\cite{dab00}, following recent interest in analyzing the gravitational lensing phenomenon in strong field regimes \\cite{vir98-vir00-tor98ab}. \\\\  However, as far as we know, literature lacks the study of an accretion process, onto a boson star, either massive or supermassive. Do the emissivity properties of the disc differ when the central object, instead of being a black hole, is assumed to be a boson star? Can we detect these differences? From where these deviations, if any, come from? Can accretion onto boson stars help model galactic centers? Of what kind? To completely answer these questions would require the analysis of different models of accretion discs, with various degrees of complexities. In this first approach, we shall analyze the properties of the simplest, steady, geometrically thin, and optically thick accretion disc model, rotating onto a static boson star. We shall compare all our results with those obtained using a black hole of the same mass.\\\\  The fact that all circular orbits are stable for static stars (as we shall show), can pose a problem to accretion scenarios upon non-baryonic (with no surface) static objects. Accretion would follow a series of stable circular orbits, loosing angular momentum and radiating part of the generated heat. If particles can always found a stable orbit, provided an enough amount of time, they would all end up in the center, and should a way of diverting them from there not exist, we would confront the formation of a baryonic black hole in the center of every non-baryonic star subject to overdense environments. This problem have apparently (as far as we are aware \\footnote{We acknowledge e-mail discussions with Dr. Tkachev in early 2000 on issues related to this point.}) been not clearly mentioned ever before in the literature of boson star solutions (see \\cite{GALAXY}), although we have no other option than confront it if we are to talk about any astrophysically relevant use of these scalar star models.  \\\\  This problem can, however, be alleviated in more general situations. In the rotating case, particles orbits were analyzed by Ryan \\cite{ryan} (see his section IV). He has shown that circular geodesic orbits are not stable beyond a given point, located at about 5/3 times the radius of the doughnut hole which appears in rotating boson star solutions. He has considered the swirling of a stellar size object (a black hole, or neutron star) within a supermassive rotating boson star and studied the gravitational wave emission as a mechanism for detection. Accretion would not continue in circular orbits since there is none after that point. Also, we are just analyzing the case in which a particle continues to travel in geodesics within the star interior, disregarding any possible influence of the boson star matter. And there is also the fact that two body encounters will be unavoidable in the innermost regions of the boson star, due to the increased matter density. These two body encounters will shift the particles to superior orbits where they'll find turning points, and bounce.\\\\  But even if an inner black hole forms, the influence upon the accreting matter that it would exert would be limited by the amount of mass it has compared with the mass of the non-baryonic object (also note that the boson star radius and the shells where most of the non-baryonic matter is located, are farther away than the Schwarzschild radius of the presumed black hole by a factor of at least 100). In general, there will be situations in which the accretion rate will be so low, that even if a black hole is formed with all the mass accreted during the lifetime of the universe, it will still have several orders of magnitude less than the boson star mass. Consider the center of the galaxy, its mass is believed to be above 2 10$^6 M_\\odot$, while the accretion rate is $\\sim 10^{-6} M_\\odot$ yr$^{-1}$. Then, in the absolutely worse case, if a black hole is formed with the accreted mass in a period of 10$^{10}$ yr, it will have 10$^4 M_\\odot$, and its gravitational influence will be defied by the non-baryonic object. Boson stars containing fermion objects within have been considered in the past \\cite{BOS-FERMION}, and this appear to be an extreme case where the fermion (neutrons) component have collapsed to a black hole. \\\\  Finally, a detailed analysis of the evolution of boson stars subject to continuous inflow of non-baryonic particles was carried out in Ref. \\cite{Suen}. Their results showed that under finite perturbations, the stars on the stable branch will settle down into a new configuration with less mass and a larger radius. Then, the accretion of non-baryonic matter possibly entering into the condensate would not pose a problem to boson star stability, nor generate a collapse. We recall that the instability of a boson star with respect to gravitational collapse has been studied by Kusmartsev et al. (see Ref. \\cite{kus91}, see also Ref. \\cite{Heusler}), using catastrophe theory. Rotating boson stars were studied (among others) by Schunck et al. \\cite{rot}.\\\\  In summary, even when the physical model used here could be regarded as too simplified and not complete (the star is not rotating, and a mechanism by which diverting matter from the center is not specifically given) we believe it still is important to begin to address the issue of real astrophysical scenarios, as accretion, upon theoretically foreseen non-baryonic objects. One of the first steps in this direction is presented in this work. ",
        "conclusions": "In this paper we have modeled a very simple accretion disc rotating around a static supermassive boson star, although we have given the scaling property that shows how to extend these results to other mass domains (equivalently, to other single boson mass cases). The disc was assumed steady, with a constant accretion rate, and thin, so that the standard theory can be applied. Throughout the paper, we have made a comparison of all results with those obtained for discs rotating around Schwarzschild black holes of the same mass. Our aim was to see whether the emissivity properties of the accretion disc are noticeably changed when the central object is. More complicated models for the accretion process as well as more realistic models for the star (as those in which the star is rotating) can be considered. We hope this work will encourage further analysis. \\\\    \\subsection*"
    },
    "0201/astro-ph0201288_arXiv.txt": {
        "abstract": "Radio relic sources in galaxy clusters are often described as the remnants of powerful radio galaxies. Here we develop a model for the evolution of such relics after the jets cease to supply energy to the lobes. This includes the treatment of a relic overpressured with respect to its gaseous surroundings even after the jets switch off. We also determine the radio emission of relics for a large variety of assumptions. We take into account the evolution of the strength of the magnetic field during the phase of relativistic particle injection into the lobes. The resulting spectra show mild steepening at around 1\\,GHz but avoid any exponential spectral cut-offs. The model calculations are used to fit the observed spectra of seven radio relics. The quality of the fits is excellent for {\\it all} models discussed. Unfortunately, this implies that it is virtually impossible to determine any of the important source parameters from the observed radio emission alone. ",
        "introduction": "Diffuse radio emission without an apparent host galaxy is found in an increasing number of clusters of galaxies \\citep[e.g.][]{gtf99,gf00,ks01,gfg01}. These objects can roughly be divided into radio halos close to the cluster centre which appear comparatively smooth in radio images and radio relics with a far more distorted and knotty appearance. The latter are usually found further away from the central regions \\citep{sr84}. Here we concentrate on radio relic sources.  A large number of possible explanations for radio relics has been proposed. All of these aim at explaining the production of a population of relativistic electrons responsible for the observed radio synchrotron emission. The suggested mechanisms range from the turbulent wakes of galaxies moving in the cluster's gravitational potential to cluster mergers. In the present paper we concentrate on the picture of relics as remnants of once powerful radio galaxies or radio-loud quasars. This scenario is based on the idea that the relics are the lobes of former radio galaxies, the jets of which ceased to supply them with energy some time ago \\citep[e.g.][]{kg94,srm01}. The relativistic electrons, the emission of which we observe today, were accelerated at the strong shocks at the ends of the active jets. The electrons lose their energy due to synchrotron losses, inverse Compton scattering of cosmic microwave background photons and, possibly, further adiabatic expansion of the relics themselves. As the relativistic electrons are not replenished by the active jets, the radiative energy losses introduce a spectral cut-off moving towards low frequencies in time. This can potentially explain the observed steep radio spectra of relics.  Note however that not all relics may be consistent with the picture of the passively fading remnant of a radio galaxy. Some relics are extended and show a rather smooth structure resembling in some ways the appearance of radio haloes, while others are composed of a number of individual knots. The variety of morphologies different from the appearance of active radio galaxies may very well indicate that our interpretation of relics as the remnants of radio galaxies is incorrect in many cases. However, it is not clear what the fate of the material forming the large scale radio structure of a formerly active radio galaxy is. Buoyant rise in the gravitational potential well of the cluster will certainly play a role \\citep[e.g.][]{cbkbf00}, but the movement of the parent galaxy through the cluster gas and cluster mergers may fragment the remnant as well. It is therefore not surprising that the morphology of relics is seldomly reminiscent of active radio galaxies. Even if all relics were indeed, as we assume here, the remnants of radio galaxies, the evolution of the remnant may imply significant departures from the simple pressure evolution envisaged in the model introduced here.  Another problem is the possibility of re-acceleration of the relativistic electrons even after the jets have switched off. \\citet{srm01} present radio observations of the relics in A13, A85, A133 and A4038. The substantial substructure in these sources may well be the signature of some highly dynamic processes being at work which could lead to further acceleration of relativistic particles. In this paper we will neglect any re-acceleration processes which may take place in relics.  The question then is whether or not a passively evolving population of relativistic electrons produces a radio spectrum consistent with observations. The main problem is that the radio spectra of relics are shallow at around 100\\,MHz and considerably steeper in the GHz; but they do not show exponential cut-offs. The spectra are therefore not consistent with a single population of electrons accelerated all at the same time and passively losing their energy. Taking this extended injection period into account, \\citet{kg94} showed that two breaks appear in the spectrum. The lower break is located at the frequency cut-off of the `oldest' electrons while the higher break is determined by the `youngest' electrons. The comparatively mild steepening of the spectra between the breaks is roughly consistent with the observations. Nevertheless, \\citet{kg94} found that for effective pitch-angle scattering of the relativistic electrons the steepening of the spectrum was still too strong. Therefore they had to suppress pitch-angle scattering to further flatten their model spectra.  The exponential cut-off in the spectrum is also avoided in the case of inhomogeneous magnetic field strengths inside the relic. The relativistic electrons spend most of their time in the low-field regions and only occasionally diffuse into the high-field regions. Thus their energy losses are reduced and by careful adjustment of the efficiency of the diffusion process an emission spectrum with mild steepening at high frequencies can be achieved \\citep{pt93,emw97}. \\citet{srm01} fitted a model based on this idea to the spectra of relic sources and found good agreement.  \\citet{kg94} assumed that the relics or their progenitors were in pressure equilibrium with their surroundings even at the time of the injection of relativistic electrons. In this paper we expand this model by taking into account the evolution of the lobes during the injection of relativistic particles. If the lobes of radio galaxies of type FRII \\citep{fr74} are the progenitors of the radio relics, then during the time the jets are still active, the lobes are expanding into the surrounding material. This implies that the pressure inside the lobes and therefore also the strength of the magnetic field are not constant during the particle injection. In fact, the strength of the magnetic field is decreasing and thus relativistic electrons injected at later times can survive for longer. We show that this effect results in radio spectra of the resulting relic sources without exponential cut-offs and a spectral slope comparable to that in observed relics.  At the time the jets switch off, the lobes may still be overpressured with respect to the ambient gas and therefore continue to expand. We derive expressions for the temporal behaviour of the volume and the pressure of such `coasting' relics. We take such a coasting phase into account in our calculation of model radio spectra from relic sources. Finally, we show that our models fit the observed spectra very well. Unfortunately, we find that none of the important source parameters, like the source age, or the physical processes taking place in the relics can be determined from the radio spectra alone.  In Section \\ref{sec:spherical} we derive the temporal behaviour of a coasting relic source after the jets have switched off. We use the results found to calculate the radio spectra of relic sources in Section \\ref{sec:radio}. Here we also show how the variation of the main model parameters influences the results. In Section \\ref{sec:compa} we fit the observed spectra of relic sources with our model and show that very little useful information about the conditions in the sources and their environments can be gained from the models. We summarise our results in Section \\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc}  We have developed a model for the evolution of the lobes of radio galaxies and radio-loud quasars of type FRII after the jets of the central AGN have stopped supplying them with energy. In the case that the lobes are still overpressured with respect to their surroundings, we show that their volume is proportional to  \\begin{equation} V \\propto \\left\\{ \\begin{array}{ll} t^{\\frac{6 \\left( \\gamma _{\\rm c}+1 \\right)}{\\gamma _{\\rm c} \\left( 7 + 3 \\gamma _{\\rm c} -2 \\beta \\right)}} & {\\rm \\ ; for \\ } \\gamma _{\\rm c} = \\gamma _{\\rm s},\\\\  t^{\\frac{6}{2 - \\beta + 3 \\gamma _{\\rm c}}} & {\\rm \\ ; for \\ } \\gamma _{\\rm c} = 4/3 {\\rm \\ and \\ } \\gamma _{\\rm s} = 5/3.\\\\ \\end{array} \\right. \\end{equation}  \\noindent The proportionality given above for $\\gamma _{\\rm c} = \\gamma _{\\rm s}$ is, strictly speaking, not a solution of the equations governing the hydrodynamics. However, we have shown that is provides an excellent fit for the numerically found correct solution. The numerical calculations also show that the transition from the active phase with jets to the coasting phase without jets is fast.  We identify the lobes in the coasting phase with radio relic sources. The calculated model spectra depend only on a small number of source and environment parameters: The total age of the source, $t$, the age of the source when the jets were switched off, $t_{\\rm s}$, and the pressure inside the lobes at that time, $p_{\\rm s}$. We consider three limiting cases for the lobe evolution: The lobes continue to expand as if the jets were still supplying energy (Model A), the lobes continue to expand but with the modified dynamics found for the coasting phase and summarised above (Model B) and the lobes come into pressure equilibrium at the time the jets switch off and stop expanding (Model C). We also investigate the effects of efficient pitch-angle scattering (JP-type models) and its absence (KP-type models).  All our models can provide satisfactory to excellent fits to the observed spectra of radio relics. Taking into account the evolution of the strength of the magnetic field inside the lobe during the injection of relativistic electrons flattens the spectra sufficiently to explain the observations. Thus we cannot rule out efficient pitch-angle scattering as suggested by \\citet{kg94}. Unfortunately, the good fits provided by all models also imply that the radio spectra of relic sources alone do not constrain any of the important source or environment parameters.  The situation is aggravated by the findings of \\citet{srm01} that the spectra can also be fitted using the alternative assumption of inhomogeneous magnetic fields inside the relics. Furthermore, even the inactive lobes may be re-ignited by compression during cluster mergers \\citep{eg01}. Note here that the model of \\citet{eg01} for radio relics does not require the radio plasma to originate in a powerful FRII-type radio galaxy.  The morphologies of radio relics are very varied and only rarely resemble the well-defined lobes of active radio galaxies. This strongly suggests that fluid flows in the clusters significantly influence the evolution of radio galaxy remnants. Any re-acceleration or energisation due to compression would make it even more difficult to determine the original conditions in the relic. A possible solution could be the combination of radio spectra with X-ray observations of sufficient spatial resolution to identify inverse Compton scattered photons originating in the relativistic plasma of the relics. The X-ray observations could then be used to estimate the electron density while the radio emission would constrain the strength of the magnetic field. Despite the problems, it is interesting to note that our simple models can fit the radio spectra of all relic sources studied here. This includes the clearly curved spectrum of Cul 0038-096, a relic with substantial substructure, as well as the spectrum of 1253+275 without any significant spectral break and a comparatively smooth appearance. This would imply that, whatever the history of the relic source, even a rather simplistic model can explain the observed spectra in a variety of different ways thus making it impossible to constrain the conditions within the relic sources. We cannot rule out a complex history or re-acceleration of relativistic electrons influencing the properties of the relic sources. However, in the light of our results, such more complicated models are not required to fit the available radio data.  The degeneracy of model parameters can, however, also be turned to advantage in the case of statistical studies of the radio source population. For example, we can use the known radio luminosity function of active FRII objects and make predictions for the cosmological distribution of their remnants. Because of the mentioned degeneracy, these predictions will be almost model-independent. Such statistical studies can be used to predict the detection rate of radio galaxy remnants in low frequency radio surveys with high surface brightness sensitivity (Cotter \\& Kaiser, in preparation).  In fact, it is not yet clear whether radio relic sources are actually the remains of the lobes of powerful radio galaxies. Deep observations of radio galaxies in which the jet flows have stopped comparatively recently are necessary to decide whether the conditions inside their lobes point towards them evolving into relics \\citep[e.g.][]{ksr00}. Clearly, the radio relic sources continue to challenge our understanding of the physics of radio galaxies and clusters."
    },
    "0201/astro-ph0201077_arXiv.txt": {
        "abstract": "Mapping and monitoring observations of SiO maser sources near the Galactic center were made with the Nobeyama 45-m telescope at 43 GHz. Rectangular mapping an area of approximately $200'' \\times 100''$ in a 30$''$ grid, and triangular mapping in a 20$''$ grid toward the Galactic center, resulted in 15 detections of SiO sources; the positions of the sources were obtained with errors of 5--10$''$, except for a few weak sources. Three-year monitoring observations found that the component at $V_{\\rm lsr}=-27$ km s$^{-1}$ of IRS 10 EE flared to about 1.5 Jy during 2000 March--May, which was a factor of more than 5 brighter than its normal intensity. Using the radial velocities and positions of the SiO sources, we identified 5 which are counterparts of  previously observed OH 1612 MHz sources. The other 10 SiO sources have no OH counterparts, but two were previously detected with VLA, and four are located close to the positions of large-amplitude variables observed at near-infrared wavelengths. A least-squares fit to a plot of velocities versus Galactic longitudes gives a rather high speed for the rotation of the star cluster around the Galactic center. The observed radial-velocity dispersion is roughly consistent with a value obtained before. It was found that all of the SiO sources with OH 1612 MHz counterparts have periods of light variation longer than 450 days, while SiO sources without OH masers often have periods shorter than 450 days. This fact suggests that lower-mass AGB stars are more often detected in SiO masers than in the OH 1612 MHz line. ",
        "introduction": "The  Galactic-center star cluster consists of mixed stellar populations (\\cite{kra95}; \\cite{mor96}). It involves a number of late-type stars which are potential candidates for OH/SiO maser emitters. Deep surveys in OH 1612 MHz, H$_{2}$O 22 GHz and SiO 43 GHz masers (\\cite{ sjo98b}b; \\cite{men97}; \\cite{izu98}) have been made; a dozen sources have been detected within a few parsec of the Galactic center. The accurate H$_{2}$O/SiO maser positions of these sources observed by the Very Large Array (VLA) provided a precise alignment between the near-infrared and radio coordinate frames (\\cite{men97}), enabling the position of Sgr A*  to be pinpointed with an accuracy better than 0.$''$1 on near-infrared images. Detections of accelerating motions of stars near the Galactic center fixed the mass of the cental object at about $3 \\times 10^{6} M_{\\odot}$ (\\cite{ghe00}). Because radio interferometers have a potential of measuring the proper motions of stars relative to Sgr A* [e.g.,  Sjouwerman et al.  (1998a)] more accurately than the present optical/infrared telescopes (\\cite{gen96}; \\cite{eck97}), to detect more SiO maser sources near the Galactic center and to investigate their properties will be quite important.  Using the  45-m telescope at Nobeyama, \\citet{izu98} detected 14 SiO sources around the Galactic center. These observations (by a beam width of about 40$''$) proved that the SiO maser source density is peaked at the Galactic center. Because this was a set of pointed observations toward the Galactic center and two one-beam offset positions, the SiO source positions were not determined with an accuracy better than the beam size. To remedy the positional uncertainties to some degree,  we made new mapping observations of SiO masers near the Galactic center using the 45-m telescope in the year 2000. The present mapping observations on a 30$''$ grid can be used to derive the source positions with an uncertainty of about 5--10$''$, depending on the signal-to-noise ratio. Also, because the intensities of SiO masers are expected to vary strongly on a time scale of one year, we monitored the intensities of SiO masers toward the Galactic center during 1999--2001. During the three-year monitoring observations, we found an SiO maser flare in the $-27$ km s$^{-1}$ component of IRS 10 EE in 2000 March--June. We present the details of these observations in this paper. ",
        "conclusions": "We have made mapping and monitoring observations of SiO maser sources near the Galactic center and have detected 15 SiO sources. Approximate positions were obtained with accuracies of about 5--10$''$; five sources were identified with the previously observed OH 1612 MHz sources. Among the sources without OH counterparts, four are close to the positions of the large-amplitude variable stars observed at near-infrared wavelengths and two to previously detected SiO sources with accurate positions from the VLA. Three-year monitoring observations of these objects found that SiO masers from IRS 10 EE flared up by a factor of more than 5 during March--May 2000. A least-squares linear fit of the velocities to the Galactic longitude in the longitude--velocity diagram gives a  high rotational speed for the star cluster around the Galactic center.  The authors thank I. Glass for stimulating discussions and reading the manuscript. They also thank M. Morris, H. Izumiura, H. Imai, and A. Miyazaki for comments. This research was partly supported by Scientific Research Grant (C2) 12640243 of Japan Society for Promotion of Sciences."
    },
    "0201/astro-ph0201307_arXiv.txt": {
        "abstract": "{ We report on two {\\it Beppo}SAX observations of BL \\,Lac (2200+420) performed respectively in June and December 1999, as part of a ToO program to monitor blazars in high states of activity. During both runs the source has been detected up to 100 keV, but it showed quite different spectra: in June it was concave with a very hard component above 5-6 keV ($\\alpha_1 \\sim 1.6$; $\\alpha_2 \\sim 0.15$); in December it was well fitted by a single power law ($\\alpha \\sim 0.6$). During the first {\\it Beppo}SAX observation BL \\,Lac showed an astonishing variability episode: the $0.3 - 2$ keV flux doubled in $\\sim 20$ minutes, while the flux above 4 keV was almost contant. This frequency--dependent event is one of the shortest ever recordered for BL\\,Lac objects and places lower limits on the dimension and magnetic field of the emitting region and on the energy of the synchrotron radiating electrons. A similar but less extreme behaviour is detected also in optical light curves, that display non-simultaneous, smaller fluctuations of $\\sim 20 \\%$ in 20 min. We fit the spectral energy distributions with a homogeneous, one-zone model to constrain the emission region in a very simple but effective SSC + external Compton scenario, highlighting the importance of the location of the emitting region with respect to the Broad Line Region and the relative spectral shape dependence. We compare our data with historical radio to $\\gamma$-ray Spectral Energy Distributions. ",
        "introduction": "Blazar objects are highly variable sources characterised by non--thermal emission that dominates from the radio to the $\\gamma$-rays. This emission is supposedly due to a relativistic jet seen at a small angle to the line of sight (Blandford \\& Rees 1978). Although blazars emit over the entire electromagnetic spectrum, their variability seems to be more pronounced in the optical-X-ray band than the radio-infrared one, both in term of flux and time scale (Ulrich et al. 1997). It is also well known that the overall spectral energy distribution (SED) of blazars shows (in a $\\nu$ vs $\\nu F_{\\nu}$ representation) two broad emission peaks; a lower frequency peak believed to be produced by synchrotron emission and a higher frequency peak probably due to the inverse Compton process. However, during strong variability events, the overall SED can change significantly. For instance, changes up to few orders of magnitude in the position of the synchrotron peak have been detected during flares in \\object{Mkn\\,501} (Pian et al. 1998) and \\object{1ES\\,2344+514} (Giommi et al. 2000). Incidentally, both sources belong to the HBL class, High Frequency Peaked blazars, with the synchrotron peak in the UV -- X-ray band. A key to understand blazar variability is the acquisition of wide band spectra during major flaring episodes. Spectral and temporal information greatly constrain the jet physics, since different models predict different variability as a function of wavelength.  To this end, we started a project to observe with the {\\it Beppo}SAX satellite (Boella et al. 1997a) blazars while they were in an active state as detected in the optical, X-ray or TeV bands. As part of this program, we already observed the two sources \\object{ON\\,231} and \\object{PKS\\,2005-489} (Tagliaferri et al. 2000, 2001). Here we present two {\\it Beppo}SAX observations of the third observed source, \\object{BL\\,Lac} itself, carried out in June and December 1999, together with simultaneous optical and radio data. Other two observations have been carried out on the objects \\object{OQ\\,530} and \\object{S5\\,0716+714} (paper in preparation).  BL Lacertae, being the prototype of the BL \\,Lac class, is one of the best--studied objects. It is a featureless LBL (low frequency peaked blazars) object, but sporadically it shows optical emission lines (EW $\\sim 6$ \\AA; Vermeulen et al. 1995; Corbett et al. 2000). In the X-ray it has been observed by many satellites; {\\it Einstein} detected a photon index of $1.68 \\pm 0.18$ (Bregman et al. 1990), Ginga a photon index in the range 1.7-2.2 (Kawai et al. 1991); ROSAT a photon index of $1.95 \\pm 0.45$ (Urry et al. 1996). ASCA observed BL\\,Lac in 1995 detecting a photon index of $1.94 \\pm 0.04$ (Madejski et al. 1999; Sambruna et al. 1999). In July 1997 following an optical outburst, it was observed with EGRET, {\\it Rossi}XTE and ASCA. EGRET found that the flux level above 100 MeV was 3.5 times higher than that observed in 1995 (Bloom et al. 1997). {\\it Rossi}XTE found a harder spectrum with a photon index in the range 1.4-1.6 over a time span of 7 days (Madejski et al. 1999). A fit to simultaneous ASCA and {\\it Rossi}XTE data shows the existence of a very steep and varying soft component below 1 keV, photon index in the range 3-5, in addition to the hard power law component with a  photon index of 1.2-1.4. Two rapid flares with time scales of 2-3 hours were detected by ASCA but only in the soft part of the spectrum (Tanihata et al. 2000). Finally, in November 1997, BL\\,Lac was observed with {\\it Beppo}SAX that detected a photon index of $1.89 \\pm 0.12$ (Padovani et al. 2001).  The {\\it Beppo}SAX observations presented here were triggered when the source was in very high optical states (R $\\sim 12.5$), but when the source was actually observed by {\\it Beppo}SAX the optical flux was lower (R = 13.4 - 13.6). In both observations the source was clearly detected up to 100 keV giving us the possibility to study BL\\,Lac over an unprecedentedly large spectral range (0.3-100 keV) with simultaneous data. ",
        "conclusions": "{\\it Beppo}SAX observations of BL \\,Lac in 1999 reveal the extremely complex behaviour of this source, as confirmed by  the comparison with previous multiwavelength observations. During the June observation, soft X-ray light curves were characterized by the fastest variability episode ever recordered for BL \\,Lac: the flux below 4 keV doubled in 20 minutes while remaining constant at higher energies. Such an amazing event allows us to put severe constraints on the dimension of the X-ray emission region, on the magnetic field and on the emitting particle energies. The frequency dependence is easily explained when performing the spectral analysis, that highlights a spectral  break attributed to the transition from the more variable synchrotron to a very hard inverse Compton spectrum: synchrotron X-ray emitting electrons are more energetic than Compton ones, so they cool faster. During the second 1999 run, {\\it Beppo}SAX detected a softer Compton component along its whole spectral range, which accounts for the constance of the light curves. The comparison of 1995 and 1997 BL \\,Lac Spectral Energy Distributions, extending to $\\gamma$-ray energies, suggests the implication of different inverse Compton emission models. Furthermore, the constance of external radiation, as inferred from the constance of emission lines, implies that the emitting shell should be differently located  with respect to the Broad Line Region, in order to explain the different observed high energy spectra. Unfortunately our observations lack simultaneous $\\gamma$-ray data, which are essential to discriminate between various high energy emission models. Therefore, we have calculated the Spectral Energy Distribution resulting from different possible locations of the emitting shell, adopting  an internal shock model. Taking into consideration the hints  given by variability and X-ray spectral shape, the most viable scenario to explain our multiwavelength BL \\,Lac observation of  1999  June implies that the emitting  shell is internal to the Broad Line Region, thus producing $\\gamma$-ray emission via Compton scattering with external photons. The 1999 December data are less revealing and we cannot discriminate between the two possible locations."
    },
    "0201/astro-ph0201131_arXiv.txt": {
        "abstract": "Previous treatments of ambipolar diffusion in star-forming molecular clouds do not consider the effects of fluctuations in the fluid fields about their mean values. This paper generalizes the ambipolar diffusion problem in molecular cloud layers to include such fluctuations. Because magnetic diffusion is a nonlinear process, fluctuations can lead to an enhancement of the ambipolar diffusion rate. In addition, the stochastic nature of the process makes the ambipolar diffusion time take on a distribution of different values. In this paper, we focus on the case of long wavelength fluctuations and find that the rate of ambipolar diffusion increases by a significant factor $\\Lambda \\sim 1 - 10$. The corresponding decrease in the magnetic diffusion time helps make ambipolar diffusion more consistent with observations. ",
        "introduction": "\\label{sec:intro}  In the usual paradigm of low mass star formation, molecular cloud cores are supported by magnetic fields. In order for star formation to take place, the cores must lose magnetic support, and this loss of support is generally thought to take place through the action of ambipolar diffusion (Mouschovias 1976; Shu 1983; Nakano 1984; Shu, Adams, \\& Lizano 1987; Lizano \\& Shu 1989; Ciolek \\& Basu 2000, 2001). This general picture has support from observations, which suggest that ion-neutral drift does indeed occur in magnetized star-forming cores (e.g., Greaves \\& Holland 1999).  An important issue facing this standard scenario is the time scale required for magnetic support to be removed from the cloud cores. As the observational picture comes into sharper focus, the number of observed cores without stars (e.g., Jijina, Myers, \\& Adams 1999) seems to be smaller than that predicted by most previous estimates from ambipolar diffusion (e.g., Ciolek \\& Mouschovias 1994, 1995; Lizano \\& Shu 1989) by a factor of 3 -- 10. In other words, loss of magnetic support by diffusion appears to be too slow, with a time scale a factor of 3 -- 10 times longer than suggested by the observed statistics of cloud cores. However, these previous calculations neglected a dimensionless factor that depends on the mass to flux ratio of the cores (Ciolek \\& Basu 2001). If the cloud cores have mass to flux ratios that approach the critical value, then the ambipolar diffusion time scale is significantly shorter than previous estimates. In particular, if the mass to flux ratio becomes supercritical, then the ambipolar diffusion time scale approaches zero. The need for this correction is bolstered by a recent compilation of Zeeman measurements of magnetic field strengths (Crutcher 1999), which suggests that many cores may have mass to flux ratios near the supercritical value. This observed sample includes only 27 cores with relatively large masses; additional measurements are necessary to clarify the observational picture.  In this study, we consider the effects of fluctuations on the mechanism of ambipolar diffusion described above. The time scale issue remains important and this work shows that ambipolar diffusion can operate more quickly in the presence of such fluctuations. In addition, because of the chaotic nature of the fluctuations, the ambipolar diffusion time scale will take on a full distribution of values for effectively ``the same'' initial states.  Fluctuations are expected to be present in essentially all star forming regions. Molecular clouds are observed to have substantial non-thermal contributions to the observed molecular line-widths (e.g., Larson 1981; Myers, Ladd, \\& Fuller 1991; Myers \\& Gammie 1999). These non-thermal motions are generally interpreted as arising from MHD turbulence (e.g., Arons \\& Max 1975; Gammie \\& Ostriker 1996; for further evidence that the observed linewidths are magnetic in origin, see Mouschovias \\& Psaltis 1995).  Indeed, the size of these non-thermal motions, as indicated by the observed line-widths, are consistent with the magnitude of the Alfv{\\'e}n speed (e.g., Myers \\& Goodman 1988; Crutcher 1998, 1999; McKee \\& Zweibel 1995; Fatuzzo \\& Adams 1993). As a result, the fluctuations are often comparable in magnitude to the mean values of the fields (T. Troland, private communication).  Background fluctuations can lead to a net change in the diffusion rate because magnetic diffusion is a nonlinear process. As many authors have derived previously (e.g., see the textbook treatment of Shu 1992), and as we present below, the (dimensionless) diffusion equation takes the schematic form \\be {\\partial b \\over \\partial \\tau} = {\\partial \\over \\partial \\mu} \\Bigl( b^2 {\\partial b \\over \\partial \\mu} \\Bigr) \\, , \\label{eq:magone} \\ee where $b$ is the magnetic field strength, $\\mu$ is the Lagrangian mass coordinate, and we have ignored density variations.  Now suppose that the magnetic field fluctuates about its mean value on a time scale that is short compared to the diffusion time (the time required for the mean value to change). We thus let $b \\to b(1+\\xi)$, where $\\xi$ is the relative fluctuation amplitude. In the simplest case in which the fluctuations are spatially independent, the right hand side of equation [\\ref{eq:magone}] is thus multiplied by a cubic factor $(1 + \\xi)^3$. Although a linear correction would average out over time, this nonlinear term must always average to a value greater than unity and the corresponding diffusion time scale grows shorter by the same factor. As an example, suppose the field spends half of its time at a value of twice its mean strength and the other half of its time near zero strength. Half of the time, the effective diffusion constant is thus larger by a factor 8, whereas the other half of the time, the diffusion constant is effectively zero. In this naive example, the mean diffusion constant is thus 4 times larger due to the fluctuations. The goal of this paper is to derive a more rigorous argument for this time scale enhancement.  This effect -- changing diffusion time scales because of fluctuations -- is well known in mathematical subfields. Simpler problems in which random noise fields drive physical systems at different rates appear in a host of textbooks (e.g., Srinivasan \\& Vasudevan 1971; Soong 1973).  More recently, in the context of ``stochastic ratchets'', it has been shown that random fluctuations can drive a physical system to propagate ``uphill'', i.e., the opposite direction of its natural propagation in the absence of fluctuations (Doering, Horsthemke, \\& Riordan 1994). In astrophysics, stochastic aspects of magnetic field fluctuations have been considered in the context of cosmic ray propagation (e.g., Jokipii \\& Parker 1969) and also in stellar atmospheres (e.g., Shore \\& Adelman 1976). Several previous papers have studied turbulent fluctuations in magnetically supported clouds, often by considering how the turbulence itself leads to field evolution in the absence of ambipolar diffusion (e.g., Kim 1997). The formation of cores through the dissipation of turbulence has been suggested (Myers \\& Lazarian 1998). Turbulence can also enhance the rate of ambipolar drift and may help explain the observed relationship between density and magnetic field strength, $B \\propto \\rho^\\kappa$ (see Zweibel 2001). In this current work, we consider ambipolar diffusion to be the main process that forms molecular cloud cores and study how fluctuations alter its rate.  This paper is organized as follows. In \\S 2, we reformulate the ambipolar diffusion calculation in a plane geometry, where we explicitly include fluctuations in both the magnetic field and the density field. We perform an analysis of the resulting set of equations in \\S 3. We specialize to the limit of long wavelength fluctuations and apply the resulting formalism to astrophysical systems. We conclude, in \\S 4, with a summary and discussion of our results. The case of short wavelength fluctuations is presented briefly in an Appendix. The most important outcome of this study is to demonstrate that fluctuations can lead to more rapid diffusion of magnetic fields in star forming regions and that the diffusion time scale takes on a distribution of values (rather than a single time scale). ",
        "conclusions": "\\label{sec:discuss}  In this paper, we have explored how fluctuations in the background fields affect the rate of ambipolar diffusion. These fluctuations force the magnetic field strength and the density to sample a distribution of values, rather than take on a single value at a given point in space and time. We have used a one-dimensional molecular cloud layer as a test problem to study the effects of such fluctuations.  The first principal result of this paper is that the time scale for ambipolar diffusion is altered by these fluctuations. In particular, fluctuations drive ambipolar diffusion to take place more rapidly, with the time scale shorter by a factor $\\Lambda \\sim 1 - 10$. For typical conditions in molecular clouds cores, the enhancement factor is near the lower end of this range, $\\Lambda \\sim 2 - 3$, but much larger enhancements remain possible.  The ambipolar diffusion time scale depends on the distribution of fluctuations. For the case of uniform distributions, for example, the ambipolar diffusion time scale varies with the amplitude $A$ as shown in Figure 1. In general, the time scale also depends on the shape of the distributions.  For a given distribution of fluctuations, the ambipolar diffusion time scale also varies from realization to realization (see Figure 2). Descriptions of the ambipolar diffusion process thus face an interesting complication, which is our second principal result: The time scale for loss of magnetic support takes on a {\\it distribution} of values instead of a single value.  Consider two identical molecular cloud regions and suppose they are laced with fluctuations following a given distribution. Because the two regions experience incomplete (and different) samplings of the fluctuation distributions as their magnetic fields diffuse outwards, they will not exhibit the same diffusion time.  This feature is more general than its manifestation in this particular test problem.  When a physical system contains an effectively random element -- in this context through chaos and turbulence -- the outcomes must be described in terms of a probability distribution. For our test problem (ambipolar diffusion in a cloud layer), the single value of the e-folding time $\\tau_{\\rm e}$ is replaced by a distribution of values (see Figure 3). Furthermore, the distribution of possible time scales approaches a gaussian form; the most likely value for the time scale is shorter than the case without fluctuations by the enhancement factor $\\Lambda = D/K$ (see equations [\\ref{eq:kdef}, \\ref{eq:finaldiff}]) and the width of the distribution is given by equation [\\ref{eq:width}].  This effect on the ambipolar diffusion time scale has important implications for star formation in molecular clouds. These clouds appear to be supported by magnetic fields and the observed magnetic field strengths are commensurate with this view. However, statistics of molecular cloud cores (with and without young stellar objects) argues that the (uncorrected) time scale for ambipolar diffusion may be too long to account for the observations. This work shows that magnetic fields can diffuse more rapidly than previous estimates suggest. This speed-up, along with any other enhancements (e.g., Ciolek \\& Basu 2001), can help account for the observed statistics of molecular cloud cores.  Another complicating issue arises: Because the ambipolar diffusion time scale takes on a distribution of values, and this distribution can be rather wide if the fluctuations change on long time scales $\\tau_X$, some core regions will experience much faster diffusion rates than others even if they have ``the same'' starting conditions.  In this regime of diffusion activity, the cores that actually form stars are those which evolve on the ``fast'' side of the distribution, whereas the cores that happen to live on the ``slow'' side of the distribution will fail to form new stars.  This preliminary treatment of fluctuations, including their effects on ambipolar diffusion and star formation, remains incomplete in several respects. In this paper, we have separated the calculation of the diffusion process from the determination of the fluctuations. In particular, we have assumed {\\it a priori} forms for the fluctuations to study their implications. In a complete treatment, one should calculate the fluctuations and their effects in a self-consistent manner. In addition, we have focused on long wavelength fluctuations and have not considered spatial gradients in the fluctuating part of the fields. Magnetic turbulence cascades down to small scales, however, so it is possible that fields fluctuate at length scales smaller than our MHD condition. This complication should also be considered in future work. Our present treatment is limited to one-dimensional slab models so that magnetic tension is not included; two-dimensional simulations should be done in the future.  Another classical problem is the heating of molecular cloud regions by ambipolar diffusion; the effects of fluctuations on this mechanism should be considered. Finally, the act of star formation provides a source of new turbulence, which drives new fluctuations and can affect the ambipolar diffusion rates of neighboring cores; this feedback effect should also be studied.  In any case, fluctuations in both the magnetic and density fields introduce an effectively random element into the ambipolar diffusion process, and thereby provide a rich class of new behavior for further study.  \\bigskip \\centerline{\\bf Acknowledgements}  We would like to thank Charlie Doering, Phil Myers, Steve Shore, and Frank Shu for useful discussions. We also thank the referee -- Glenn Ciolek -- for comments that improved the paper.  MF is supported by the Hauck Foundation through Xavier University. FCA is supported by NASA through a grant from the Origins of Solar Systems Program and by Univ. Michigan through the Michigan Center for Theoretical Physics.  \\bigskip \\centerline{\\bf APPENDIX A: Cylindrical Geometry} \\medskip  In this Appendix, we consider the effects of fluctuations on ambipolar diffusion in a molecular cloud filament. We consider only the simplest case of magnetic field lines that are aligned with the axis of the filament and depend only on the radial coordinate $r$, i.e., we have $$ {\\bf B} = B(r) \\, {\\hat z} \\, . \\eqno({\\rm A}1)$$ With this basic configuration, the equations of motion take the form $$ {\\partial \\rho \\over \\partial t} + {1 \\over r} {\\partial \\over \\partial r} (r \\rho u) = 0 \\, , \\eqno({\\rm A}2)$$ $$ {\\partial u\\over\\partial t} + u {\\partial u\\over \\partial r} = g - {1\\over\\rho}{\\partial\\over\\partial r} \\left(P+{B^2\\over 8\\pi}\\right)\\, , \\eqno({\\rm A}3)$$ $$ {1\\over r}{\\partial\\over\\partial r}(rg) = -4\\pi G\\rho\\,, \\eqno({\\rm A}4)$$ and finally $$ {\\partial B \\over \\partial t} + {1 \\over r} {\\partial \\over \\partial r} (r B u) =  {1 \\over 4 \\pi \\gamma C} {1 \\over r} {\\partial \\over \\partial r}  \\Bigl\\{ {r \\over \\rho^{3/2}} B^2 {\\partial B \\over \\partial r}  \\Bigr\\} \\, , \\eqno({\\rm A}5)$$ where we have defined $u$ to be the radial (and only nonvanishing) component of the velocity, and we have made use of the relationship defined by equation [\\ref{eq:rhoion}]. In practice, these filaments will be subject to clumping instabilities in both the linear (Gehman, Adams, \\& Watkins 1996) and nonlinear regimes (Adams, Fatuzzo, \\& Watkins 1994); in this derivation, however, we neglect this issue and focus on the effects of the cylindrical geometry.  Next, we introduce the fluctuations through the ansatz given by equation [\\ref{eq:ans}] and rewrite the problem in terms of a Lagrangian description of the dynamics (e.g., see \\S 2).  For the cylindrical geometry considered here, the relevant Lagrangian coordinate is the mass per unit length $\\sigma$ along the filament within a radius $r$ of the central axis, i.e., $$ \\sigma \\equiv \\int_0^r \\, \\rho(r', t) \\, r' \\, dr' \\, . \\eqno({\\rm A}6)$$ Notice that the variable $\\sigma$ differs from the true mass per unit length by a factor of $2 \\pi$ which has been omitted for simplicity.  The original problem in the variables $(r, t)$ is now transformed to one in new variables $(\\sigma, t)$ and the derivatives transform according to $$ {\\partial \\over \\partial t} + u {\\partial \\over \\partial r} \\to \\, {\\partial \\over \\partial t}\\Bigg|_\\sigma \\, , \\eqno({\\rm A}7)$$ $$ {\\partial \\over \\partial r} \\to \\, r \\rho {\\partial \\over \\partial \\sigma}\\Bigg|_t \\, , \\eqno({\\rm A}8)$$ $$ u \\to {\\partial r \\over \\partial t}\\Bigg|_\\sigma \\, , \\label{eq:concyl} \\eqno({\\rm A}9)$$  With this transformation, the equation of continuity becomes $$ {\\partial r \\over \\partial \\sigma} = {1 \\over r \\rho} \\, . \\eqno({\\rm A}10)$$ For an isothermal equation of state, the force equation can be written in the form $$ - {1 \\over r} \\, {\\partial^2 r \\over \\partial t^2} = {4\\pi G\\over r^2} \\int_0^r \\rho (1 + \\eta) r' dr' + {a^2\\over 1+\\eta} {\\partial \\over \\partial \\sigma} [\\rho(1+\\eta)] + {1 \\over 1+\\eta} {\\partial  \\over \\partial \\sigma} \\left[{B^2 (1+\\xi)^2\\over 8\\pi}\\right]\\, , \\eqno({\\rm A}11)$$ and the nonlinear diffusion equation for the magnetic field becomes $$ {\\partial \\over \\partial t} \\left[{B(1+\\xi)\\over \\rho(1+\\eta)} \\right]=  {1\\over 1+\\eta}{\\partial \\over \\partial \\sigma} \\Bigl\\{ {r^2 B^2 (1+\\xi)^2\\over 4\\pi \\gamma C \\rho^{1/2} (1+\\eta)^{3/2}}{\\partial  \\over \\partial \\sigma}  [B(1+\\xi)] \\Bigr\\} \\, . \\label{eq:difcyl} \\eqno({\\rm A}12)$$ Equations [A10 -- A12] exhibit exactly the same form as their slab counterparts (equations [\\ref{eq:continuity} -- \\ref{eq:diffuse0}]). As such, the effects of fluctuations on ambipolar diffusion in a cylindrical filament will scale exactly as for the slab geometry considered in the main text.  \\bigskip \\centerline{\\bf APPENDIX B: Short Wavelength Fluctuations} \\medskip  We consider the effects of short wavelength fluctuations in this Appendix.  As noted in the main body of the text, this analysis is complicated by the fact that the solutions to stochastic differential equations depend on the manner in which various limits are taken (Doering 1990). The formulation presented below thus represents one possible approach to the general problem.  The most likely source of short wavelength fluctuations is MHD turbulence, which is present in most regions of molecular clouds. The MHD condition, already built into the ambipolar diffusion equations, requires that the neutral fluid remain coupled to the ions and to the magnetic field.  Physically, this condition is met if the ion-neutral collision frequency $f_{in} = \\gamma \\rho_i$ exceeds the frequency associated with the MHD turbulence. The latter frequency can be approximated by $f_{mhd} \\approx v_A / \\lambda$, where $v_A = B / (4 \\pi \\rho)^{1/2}$ is the Alfv\\'en wave speed and $\\lambda$ is the length scale of the fluctuations.  As a result, the coupling condition requires that $$ \\chi > 0.09 {b\\over p}\\, , \\eqno({\\rm B}1) $$ where $\\chi$ is the dimensionless turbulence length scale as defined by equation [\\ref{eq:chidef}].  If we assume that the fluctuations are both spatially and temporally symmetric (which means that $\\xi$ and $\\eta$ are not correlated with their first order derivatives), then the following relations hold: $$ \\left< F(\\xi, \\eta) {\\partial\\eta\\over\\partial\\mu} \\right> \\approx 0 \\qquad {\\rm and} \\qquad \\left< F(\\xi, \\eta) {\\partial \\xi \\over\\partial\\mu} \\right> \\approx 0 \\, , \\eqno({\\rm B}2) \\label{eq:fslope} $$ and $$ \\left< F(\\xi, \\eta) {\\partial\\eta\\over\\partial\\tau} \\right> \\approx 0 \\qquad {\\rm and} \\qquad \\left< F(\\xi, \\eta) {\\partial \\xi \\over\\partial\\tau} \\right> \\approx 0 \\, , \\eqno({\\rm B}3) $$ for all well-behaved functions $F(\\xi, \\eta)$. Under these conditions, the quasi-equilibrium state described in \\S 3.2 remains valid.  For short wavelength fluctuations, expanding the diffusion equation using the same approach as presented in \\S 3.3 yields the form $$ (1 + \\xi) {\\partial\\over\\partial\\tau} \\Bigl( {b \\over p} \\Bigr) + {b \\over p} \\Bigl[ {\\dot \\xi} - {\\dot \\eta} \\, { (1 + \\xi) \\over (1+\\eta) } \\Bigr] = { (1+\\xi)^3 \\over (1+\\eta)^{3/2} } {\\partial \\over \\partial\\mu} \\left\\{ {b^2 \\over p^{1/2}} {\\partial b \\over \\partial\\mu} \\right\\} + \\left\\{ {\\partial\\over\\partial\\mu} \\left[{(1+\\xi)^3\\over (1+\\eta)^{3/2}}\\right] \\right\\} \\left( {b^2 \\over p^{1/2}} {\\partial b\\over\\partial\\mu}\\right) $$ $$ +{(1+\\xi)^2 \\over (1+\\eta)^{3/2}} \\left\\{ {\\partial\\over\\partial\\mu} (1+\\xi) \\right\\} \\left\\{ {\\partial\\over\\partial\\mu} \\left[{b^3\\over p^{1/2}}\\right]\\right\\} +{b^3 \\over p^{1/2}} {\\partial\\over\\partial\\mu} \\left\\{{(1+\\xi)^2\\over (1+\\eta)^{3/2}} {\\partial\\over\\partial\\mu} (1+\\xi)\\right\\} \\, . \\eqno({\\rm B}4) \\label{eq:bigdiff} $$ We note that the second and third terms on the right hand side vanish when they are time averaged because they take the form given by equation [\\ref{eq:fslope}]. With this simplification, the time-averaged diffusion equation reduces to the form $$ {\\partial\\over\\partial\\tau} \\left({b\\over p}\\right) = D {\\partial\\over\\partial\\mu}\\left( {b^2\\over p^{1/2}} {\\partial b\\over\\partial\\mu}\\right) + G {b^3\\over p^{1/2}}\\,, \\eqno({\\rm B}5) \\label{eq:fdiffsw} $$ where $D$ is defined in equation [\\ref{eq:finaldiff}] and $$ G = \\left<{\\partial\\over\\partial\\mu} \\left[{(1+\\xi)^2\\over (1+\\eta)^{3/2}} {\\partial\\over\\partial\\mu} (1+\\xi)\\right]\\right>\\,. \\eqno({\\rm B}6) $$ For the case in which the fluctuations $\\xi$ and $\\eta$ are uncorrelated, the parameter $G$ simplifies to the form $$ G = {1\\over 3}\\left<{1\\over (1+\\eta)^{3/2}}\\right> \\left<{\\partial^2 \\over \\partial\\mu^2} (1+\\xi)^3\\right> \\, . \\eqno({\\rm B}7) $$ Similarly, for the case of perfectly correlated fluctuations, $G$ simplifies to the form $$ G = {2\\over 3} \\left<{\\partial^2 \\over \\partial\\mu^2} (1+\\xi)^{3/2}\\right> \\, . \\eqno({\\rm B}8) $$  The relative size of the two terms on the right hand side of equation [B5] ultimately determines the behavior of the magnetic field diffusion. Since the fluctuations $\\xi$ and $\\eta$ vary on a length scale $\\chi \\ll 1$, whereas $b$ and $p$ vary on a lengthscale $1 + \\alpha_0 \\approx K v_A^2 / a^2 \\gg 1$, a simple scaling analysis naively suggests that the second term (with coefficient $G$) would dominate over the first (with coefficient $D$). Upon closer inspection, however, we see that the derivatives of the fluctuations tend to cancel out, so that the relative sizes of $D$ and $G$ depend on the form of the fluctuations. In any case, however, this treatment does not yield an expression that can be scaled to the previous solutions with no fluctuations.  \\newpage"
    },
    "0201/astro-ph0201495_arXiv.txt": {
        "abstract": "To understand luminous AGNs in the $z<1$ universe, the {\\it ASCA} AGN samples are the best at present. Combining the identified sample of AGNs from {\\it ASCA} Large Sky Survey and Medium Sensitivity Survey, the sample of hard X-ray selected AGNs have been expanded up to 108 AGNs above the flux limit of 10$^{-13}$ erg s$^{-1}$ cm$^{-2}$ in the 2--10~keV hard X-ray band. We discuss the fraction of absorbed AGNs in the hard X-ray selected AGN sample, and nature of absorbed luminous AGNs. ",
        "introduction": "\\begin{figure} \\plotone{Figure1.eps} \\caption{ Left) $R$-band magnitudes of optical counterparts of ALSS (square) and AMSSn (circle) AGNs are plotted as a function of 2--10~keV hard X-ray flux. Dashed lines represent the X-ray to optical flux ratio of $\\log f_{X}/f_{V} = +2$,$+1$,0,$-1$,and $-2$ from top to bottom. Triangles and asterisks indicate samples from {\\it HEAO1} A2 (Piccinotti et al. 1982) and {\\it Chandra} survey in HDF-N (Hornschemeier et al. 2001) Right) Hard X-ray luminosities of the hard X-ray selected AGNs plotted as a function of redshift. Same symbols as in the left panel.} \\end{figure}  The fraction of absorbed AGNs, especially luminous absorbed AGNs, is one of a big issue in understanding the true number density of active nuclei in the universe. Recently many candidates of absorbed luminous AGNs are found in AGN surveys in radio, X-ray, and near-infrared wavelengths (e.g., Webster et al. 1995). The discoveries imply that we have been missing significant fraction of nucleus with high activity in traditional optical/UV-selections of AGNs due to absorption to the nucleus. However the fraction of the absorbed AGN in the entire AGN population is not clear. Radio-selected samples are affected by red AGNs with red synchrotron component (Francis et al. 2001), soft X-ray selection is biased against heavily absorbed AGNs (Kim \\& Elvis 1999), and 2MASS-selected red AGNs are limited in the low redshift universe (Cutri et al. in this volume).  In order to construct a complete sample of AGNs less biased against absorption to nucleus, we conduct optical follow-up observations for {\\it ASCA} Large Sky Survey (hereafter ALSS; Ueda et al. 1999) and {\\it ASCA} Medium Sensitivity Survey (hereafter AMSS; Ueda et al. 2001) in the hard X-ray band. Hard X-ray emission can penetrate the obscuring matter of absorbed AGNs and is very suitable in searching absorbed AGNs. Using 2--10~keV hard X-ray emission, we can detect AGNs with X-ray absorption up to hydrogen column density of 10$^{22\\sim23}$ cm$^{-2}$, which corresponds to $A_{V}$ of 20 $\\sim$ 50 with galactic conversion factor, without bias. ALSS is a survey in a continuous field with 5.4 square degree near the north galactic pole. We selected 34 X-ray sources detected with SIS 2--7~keV significance larger than 3.5$\\sigma$. The sources are identified with 30 AGNs, 2 clusters of galaxies and 1 galactic star (Akiyama et al. 2000). One X-ray source with hard spectrum is still unidentified, and {\\it Chandra} follow-up observation is planed in Cycle 3. AMSS is a serendipitous source survey based on {\\it ASCA} pointing observations conducted in high galactic latitude region ($|b|>20^{\\circ}$). We conducted optical follow-up observations for 86 X-ray sources detected with GIS 2--10~keV significance larger than 5.6$\\sigma$ in the northern sky (declination above $20^{\\circ}$; we call AMSSn sample). All of the X-ray sources are identified with 78 AGNs, 7 clusters of galaxies, and 1 galactic star (Akiyama et al. in preparation). In total, we constructed sample of 108 hard X-ray selected AGNs with the flux limit of {\\it ASCA}, about $\\sim 10^{-13}$ erg s$^{-1}$ cm$^{-2}$ in the 2--10~keV band. In Figure 1, we plotted the hard X-ray flux vs. optical magnitude (left) and the redshift vs. luminosity distribution (right) diagrams of ALSS and AMSSn AGNs. The {\\it ASCA} samples are two orders of magnitude brighter and more luminous than the sample of deep {\\it Chandra} and {\\it XMM-Newton} surveys, and consists of luminous AGNs, i.e., QSOs,  in the universe below redshift 1. The high completeness of the {\\it ASCA} samples makes us possible to discuss the fraction of absorbed AGNs definitely. ",
        "conclusions": ""
    },
    "0201/astro-ph0201176_arXiv.txt": {
        "abstract": "Density waves in the central kpc of galaxies, taking the form of spirals, bars and/or lopsided density distributions, are potential actors of the redistribution of angular momentum. They may thus play an important role in the overall evolution of the central structures, not mentioning the possible link with the active/non-active central mass concentration. I present here some evidence for the presence of such structures, and discuss their importance in the context of dynamical evolution. ",
        "introduction": "Spirals, bars, warps and lopsidedness are ubiquitous at large scales in disk galaxies (e.g. Sancisi 1976; Kamphuis et al. 1991). These structures are usually understood as density waves propagating in the stellar and/or gaseous component. In this paper, I emphasize the case for density waves in the central kpc of disc galaxies, and illustrate their presence and role with a few examples.  \\subsection{Role of density waves}  Density waves such as bars, spirals and lopsided density distributions have been recognised as potential actors in the secular evolution of galaxies, partly because they are non-axisymmetric perturbations of the potential. Gas flows within a bar (Athanassoula 1992) are very illustrative of how such structures can redistribute angular momentum: gas can be transported e.g. inwards (and angular momentum outwards) via driven trailing spirals, strong shocks can occur and rings may form (usually near specific resonances). All these correspond to observed signatures of bars. Density waves also tend to heat the stellar system, therefore acting against subsequent gravitational instabilities. Small versions of the large-scale waves observed in disk galaxies have been invoked as responsible for the fueling of the central 10s of parsecs, and hypothetically the central AGN and/or starburst (Schlosman, Frank \\& Begelman 1989; Regan \\& Mulchaey 1999). But are these inner bars/spirals modes similar to their large-scale parents?  \\subsection{Towards the centre}  A dynamical probe which would travel towards the centre of a galaxy would see a number of changes occuring: spatial sizes are getting smaller and time-scales shorter (a useful number to keep in mind is: $10^6\\,$yr corresponds to $\\sim 100$~pc at 100~km.s$^{-1}$), the spheroidal component starts to significantly contribute to the potential, and eventually, at a scale of a few parcsecs, the central cusp and/or supermassive black hole may dominate the potential. Within the central kpc, an inner cold component (e.g. a disk) may be less self-gravitating, and tend to evolve more rapidly. These changes obviously affect the growth and propagation of density waves in the inner parts of galaxies. As observers better focus on the central kpc of disk galaxies (with the help of an improved spatial resolution), structures such as inner spirals, bars and lopsided density distribution are revealed (Erwin et al. 2001; Regan \\& Mulchaey 1999), and are indeed usually weaker than their large-scale versions. Before these are confirmed as true density waves however, their kinematics should be studied in more detail. After a brief reminder on resonances, I will provide a few examples of such observations.  \\subsection{Epicycle approximation}  Circular orbits can be used as zeroth order approximations for orbits in the equatorial plane of a thin disk. The first order terms can easily be evaluated by linearising the equations of motion in that plane: a retrograde epicycle motion is added on the circular orbit. This epicycle is the combination of radial and azimuthal harmonic oscillations sharing the same frequency $\\kappa$, which only depends on the first and second radial derivatives of the potential (see Dehnen 1999 for an alternative view on Lindblad's epicycle theory). The shape of the (first order) orbits at a certain radius will then depend on the ratio $\\kappa / \\Omega$ where $\\Omega$ is the circular frequency: when this ratio is an integer $m$, the orbits are periodic (closing after $m$ epicycles). With the addition of a density wave with a pattern speed $\\Omega_p$, the potential stays constant only in a frame rotating with the wave. In this rotating frame, the important quantity which decides on the shape of the orbit then becomes $\\kappa / \\left(\\Omega - \\Omega_p\\right)$. The resonances are thus located at radii where this ratio takes integer ($m$) values: $ \\Omega_p = \\Omega + \\kappa / m $. This is where the disc potential (with its natural frequency $\\kappa$) and the wave (of angular frequency $\\Omega_p$) may interact. The most important are the Lindblad Resonances (LR; the Inner LR or ILR for $m=2$, and the Outer LR, or OLR for $m=-2$) and the Corotation Resonance (CR) where $\\Omega = \\Omega_p$ ($m \\rightarrow \\infty$). At the resonances, the (linearised) orbits are closed in the frame rotating with the wave, thus defining the different families of orbits from which the skeleton of the system is built. ",
        "conclusions": "The short timescales associated to the physical processes occuring in the central regions of galaxies do not help our understanding of their morphology and dynamics. We should then view them as continuously evolving systems involving several recurrent interlinked and non simultaneous processes. Density waves are certainly playing an important role in this game. If inner bars and spirals are already on the priority list of observers and theoreticians, then $m=1$ modes must now be seriously considered as potential players on the evolutionary stage of galactic centres."
    },
    "0201/astro-ph0201340_arXiv.txt": {
        "abstract": "New estimates of the distances of 36 nearby galaxies are presented based on accurate distances of galactic Cepheids obtained by Gieren, Fouqu\\'e and Gomez (1998) from the  geometrical Barnes-Evans method.  The concept of 'sosie' is applied to extend the distance determination to extragalactic Cepheids without assuming the linearity of the PL relation. Doing so, the distance moduli are obtained in a straightforward way.  The correction for extinction is made using two photometric bands ($V$ and $I$) according to the principles introduced by Freedman and Madore (1990). Finally, the statistical bias due to the incompleteness of the sample is corrected according to the precepts introduced by Teerikorpi (1987) without introducing any free parameters (except the distance modulus itself in an iterative scheme).  The final distance moduli depend on the adopted extinction ratio ${R_V}/{R_I}$ and on the limiting apparent magnitude of the sample. A comparison with the distance moduli recently published by the Hubble Space Telescope Key Project (HSTKP) team reveals a fair agreement when the same ratio ${R_V}/{R_I}$ is used but shows a small discrepancy at large distance.  In order to bypass the uncertainty due to the metallicity effect it is suggested to consider only galaxies having nearly the same metallicity as the calibrating Cepheids (i.e. Solar metallicity). The internal uncertainty of the distances is about 0.1 magnitude but the total uncertainty may reach 0.3 magnitude. ",
        "introduction": "As an extension of our study of the kinematics of the local universe (KLUN+) we need an accurate value for the global Hubble constant and accurate distances of individual galaxies. The calibration of the distance scale is thus a fundamental step in this process. The aim of this work was to calibrate the distance scale from nearby galactic Cepheids for which the HIPPARCOS satellite measured geometrical parallaxes. This should avoid the step of calibrating the distance scale by assuming a given distance to the Large Magellanic Cloud (LMC). Unfortunatelly, it turns out that these measurements are very difficult to use due to a statistical bias (Lutz and Kelker, 1973). The difficulties can be solved by proper treatment, like the one proposed by Feast and Catchpole (1997). It has been shown that this leads to unbiased results (Pont et al., 1997; Lanoix et al. 1999),  On the other hand, individual measurements of Cepheids from HIPPARCOS are relatively inaccurate because of the distance of galactic Cepheids. Excluding $\\alpha$ UMi which does not pulsate in the fundamental mode, the best geometrical parallax of an individual Cepheid obtained from HIPPARCOS is 3.32$\\pm$ 0.58 marcsec for $\\delta$ Cephee. This leads to an uncertainty in the distance modulus of 0.38 magnitude. In comparison, the {\\it quasi-geometrical} method of Barnes-Evans applied to Cepheids (Gieren, Fouqu\\'e, Gomez, 1998; hereafter GFG), gives distance moduli with a typical uncertainty less than 0.1 magnitude (the external error can be estimated to about 0.2 magnitude according to Table 7 in GFG). We call this method {\\it quasi-geometrical} because it requires only a few assumptions. The method is independent of any determination of the LMC distance and has a relatively small systematic error (about 0.2 magnitude). {\\it Thus, we decided to calibrate the distance scale using the work done by Gieren, Fouqu\\'e and Gomez (1998).}  Nevertheless, other difficulties appear. The slope of the Period-Luminosity relation (hereafter, PL relation) determined from the adopted calibrating galactic Cepheids differs from the slope obtained for the LMC by the same authors (GFG) (Table \\ref{slopes}).. For the LMC, the slopes in V and I bands are now confirmed by the OGLE survey (Udalski. et al., 1999). What slope should we adopt?  \\begin{table} \\caption{Slopes of the PL relation.} \\label{slopes} \\begin{tabular}{lll} \\hline source & $a_V$ &  $a_I$  \\\\ \\hline GFG(MW)  & $-3.037\\pm 0.138$  & $-3.329 \\pm 0.132$  \\\\ GFG(LMC)  & $-2.769\\pm 0.073$  & $-3.041 \\pm 0.054$  \\\\ OGLE(LMC) & $-2.765$  &$-2.963$ \\\\ \\hline \\end{tabular} \\end{table}  The true physical relation is actually a Period-Luminosity-Color (hereafter, PLC) relation written as $M=\\alpha \\log P + \\beta C_o + \\gamma$, where $M$ is the absolute magnitude and $C_o$ the intrinsic color. The PL relation is simply the projection of the PLC onto the P-L plane. In the PLC relation the slope $\\partial M / \\partial logP$ is constant. However, the observed slope of the PL relation depends on the distribution of observed Cepheids in the PLC plane (i.e., on the color distribution of the sample). Hence, the slope in a given photometric band may partially depend on the metallicity, because it affects the intrinsic color. Linear non-adiabatic models do predict that the slope is constant when one uses bolometric magnitudes (Baraffe et al., private communication), whereas non-linear models predict that the slope depends on the metallicity also for the bolometric magnitudes (Bono et al., 2000 and references therein) and predict that the slope in a given band depends on the metallicity. Because the metallicity of the LMC differs from the metallicity in the Solar neighbourhood, the choice of slopes in different bands is difficult. {\\it In order to avoid this dilemma we decided to apply the method of 'sosie' (Paturel, 1984) because it does not require knowledge of the slope and zero point of the PL relation \\footnote{this method was first introduced to solve the same kind of problems for the Tully-Fisher relation (1977)}.}  The correction for extinction produced by interstellar matter is another difficulty. It can be solved by assuming that the extinction law is universal. We will thus assume that the extinction on an apparent magnitude is proportional to the color excess ($A_{\\lambda} = R_{\\lambda} (C-C_o)$, where $C$ is the reddened color). The factor of proportionality $R_{\\lambda}$ is taken from tabulations (e.g., Cardelli, Clayton \\& Mathis, 1989 ; Caldwell \\& Coulson, 1987 ; Laney \\& Stobie, 1993). It depends on both the considered band and color. With such an assumption it is possible to use the Freedman and Madore (1990) precepts of de-reddening. Two bands are needed in order to calculate a color. Because most extragalactic Cepheids are measured in V- and I-band from The {\\it Hubble Space Telescope} (hereafter, HST), we will use these two bands. {\\it Thus, the Freedman and Madore (1990) de-reddening method will be adapted to the sosie method, used in V and I photometric bands.}  Finally, an ultimate difficulty comes from the incompleteness bias. This bias was first studied by Teerikorpi (1987) for application to galaxy clusters (Bottinelli et al., 1987). It was first denounced by Sandage (1988) in application to the PL relation and re-discussed later by Lanoix, Paturel and Garnier (1999a). The sample to which we are applying the PL relation must be statistically representative of the calibrators themselves. Indeed, due to the intrinsic scatter of the PL relation, there is a given distribution of absolute magnitudes at a given period. At increasing distances the fainter end of this distribution is progressively missed and the distribution of the actual sample changes. Restricting the sample to Cepheids with a period larger than a given limiting period reduces this bias. The limiting period depends on a first estimate of the distance, on the apparent limiting magnitude and on the characteristics of the PL relation (dispersion, slope and zero-point). In fact, the full theory of Teerikorpi is applicable. The method is much more complete than the rough rule of thumb used as a quick approach in an application in which a detailed treatment was not needed. However, we want to derive final distance moduli and the precise bias correction must be used. Note that the slope and zero point of the PL relation are needed but only as second order terms and thus, the uncertainties mentioned about their choice do not present any significant difficulty (this will be confirmed in section 4.3). {\\it The incompleteness bias will be corrected using the precepts given by Teerikorpi (1987).}   In section 2 we will describe the material used for this study: the calibrating sample by GFG and our extragalactic Cepheid database (Lanoix et al., 1999b).  In section 3 we describe the 'sosie' method  and give the basic equation for the calculation of the distance modulus of an extragalactic Cepheid.  In section 4 we give the results obtained for 1840 Cepheids belonging to 36 nearby galaxies described in the previous section. We also discuss these results and compare them with those recently published by Freedman et al. (2001). ",
        "conclusions": "The distance scale can be calibrated using galactic Cepheids. LMC provides us with numerous Cepheids located at the same distance. This gives a way to derive an accurate slope for the Cepheid PL relation. But its low metallicity (with respect to most of the galaxies of the sample) is a cause of suspicion; we are not sure that this slope can be applied to all kinds of metallicity.  So, we preferred, in a first step, to calibrate the distance scale by using accurate distances of galactic Cepheids published by Gieren, Fouqu\\'e and Gomez (1998). These distances are based on the geometrical Barnes-Evans method.  Further, we applied the concept of 'sosie' (Paturel, 1984) to extend distance determinations to extragalactic Cepheids without having to know either the slope or the zero-point of the PL relation. The distance moduli are obtained in a straightforward way. For the calibrating galactic Cepheids we checked the internal coherence from the same method.   The correction for the extinction is made by using two bands ($V$ and $I$) according to the principles introduced by Freedman and Madore (1990). There is no need for color excess estimation.  Finally, the incompleteness bias is corrected according to the precepts introduced by Teerikorpi (1987). Without any free parameters (except the distance modulus itself), the bias curve calculated for each individual host galaxy fits very well the observed distance moduli. This gives us confidence in our final distance moduli. Nevertheless, the small departure from the measurements published recently by Freedman et al. (2001) at distances larger than 10Mpc ($\\mu=30$) must be clarified.  In order to bypass the uncertainty due to metallicity effects it is suggested to consider only galaxies having nearly the same metallicity as the calibrating Cepheids (i.e. Solar metallicity). In Table \\ref{param} the distance moduli that can be considered as more secure are noted with an asterisk ($\\ast$). Galaxies with $\\Delta \\mu$ larger than $\\approx 0.3$ mag. or with small $n$ do not receive this flag. For a given ratio ${R_V}/{R_I}$, the uncertainty of the distances is about 0.1 magnitude but the total uncertainty may be about 0.3 magnitude. The choice of a given ${R_V}/{R_I}$ ratio is a first source of uncertainty. The actual ratio depends on the extinction law in our Galaxy, on the extinction law in the host galaxy and on the color of the considered Cepheid. For the future it would be interesting to search for a clue allowing us to decide which value is the best in a given direction for a Cepheid in a given host galaxy. The proper determination of the limiting magnitude of the sample is a second source of uncertainty. It can be accurately determined only when a large number of Cepheids is available to provide us with good statistics.  Presently, the calibration of the distance scale can barely be better than $\\sigma_{\\mu}=0.3$ magnitude. Thus, the uncertainty on the Hubble constant, $\\sigma(H) \\approx \\sigma_{\\mu}  H / 5 \\log{e}$, cannot be better than  about $10 km.s^{-1}.Mpc^{-1}$.   \\begin{appendix}"
    },
    "0201/astro-ph0201389_arXiv.txt": {
        "abstract": "We demonstrate that the atomic alignment of the hyperfine-structure components of the ground level S$_{1/2}$ of \\ion{Na}{1} and of the upper level P$_{1/2}$ of the D$_1$ line are practically negligible for magnetic strengths $B>10\\,\\rm G$, and virtually zero for $B\\ga 100\\,\\rm G$. This occurs independently of the magnetic-field inclination on the stellar surface (in particular, also for vertical fields). Consequently, the characteristic antisymmetric linear-polarization signature of the scattered light in the D$_1$ line is practically suppressed in the presence of magnetic fields larger than 10~G, regardless of their inclination. Remarkably, we find that the scattering polarization amplitude of the D$_2$ line increases steadily with the magnetic strength, for vertical fields above 10~G, while the contribution of alignment to the polarization of the D$_1$ line rapidly decreases. Therefore, we suggest that spectropolarimetric observations of the ``quiet'' solar chromosphere showing significant linear polarization peaks in both D$_1$ and D$_2$ cannot be interpreted in terms of one-component magnetic field models, implying that the magnetic structuring of the solar chromosphere could be substantially more complex than previously thought. ",
        "introduction": "\\label{sec:intro}  In a recent work, one of the authors (\\citealt{LA98}) concluded that his explanation in terms of ground-level atomic polarization of the ``enigmatic'' linear polarization peaks of the \\ion{Na}{1} D-lines, observed by \\citet{STEN97} in ``quiet'' regions close to the solar limb, implies that the magnetic field in the lower solar chromosphere must be either isotropically distributed and extremely weak (with $B\\la 0.01\\,\\rm G$) or, alternatively, practically radially oriented. That investigation was based on a formulation of line scattering polarization that is valid in the absence of magnetic fields. The suggestion that the magnetic field of the lower solar chromosphere cannot be stronger than about 0.01 G unless it is oriented preferentially along the radial direction was based on the sizeable amount of ground-level polarization required to fit the $Q/I$ observations of \\citet{STEN97}, and on the assumption that the atomic polarization of the ground-level of \\ion{Na}{1} must be sensitive to much weaker magnetic fields than the atomic polarization of the upper levels of the D$_1$ and D$_2$ lines.  On the whole, Landi Degl'Innocenti's (\\citeyear{LA98}) argument that the observed linear polarization peaks in the cores of the \\ion{Na}{1} D-lines are due to the presence of ground-level atomic polarization seems very convincing. However, for a rigorous interpretation of spectropolarimetric observations (e.g., \\citealt{VMP01,STEN01}) it is of fundamental importance to clarify the physical origin of this polarization by carefully investigating how it is actually produced, and modified by the action of a magnetic field of given strength and inclination.\\footnote{Remarkably, some useful information can be found in the atomic physics literature, notably in the paper by \\citet{ELL34} regarding their determination of hyperfine separation constants, and in the work of \\citet{LEH69} concerning the orientation of the diamagnetic ground state of Cadmium by optical pumping.} ",
        "conclusions": "\\label{sec:conclu}  One of the most interesting results of this investigation is that the atomic polarization of the HFS levels of the ${\\rm S}_{1/2}$ and ${\\rm P}_{1/2}$ states of \\ion{Na}{1} is practically negligible for $B > 10\\,\\rm G$, and virtually vanishes for $B {\\ga} 100\\,\\rm G$, even for a purely vertical field. Consequently, the characteristic antisymmetric scattering polarization signature of the D$_1$ line is practically suppressed in the presence of fields larger than 10~G, regardless of their inclination.  Concerning the observable effects, we find that the scattering polarization amplitude of the D$_2$ line increases steadily with the magnetic strength, in the case of vertical fields larger than 10~G, whereas the contribution of atomic alignment to the linear polarization of the D$_1$ line rapidly decreases. On the contrary, for vertical fields such that $B\\la 10\\,\\rm G$ (or, alternatively, for turbulent or canopy-like fields with a predominance of much weaker fields) it is possible to have a non-negligible scattering polarization signal for the D$_1$ line, but then the maximum D$_2$ core amplitude corresponds to the $B=0\\,{\\rm G}$ case. From this we tentatively conclude that spectropolarimetric observations of the ``quiet'' solar chromosphere showing significant scattering polarization peaks in both the D$_1$ and D$_2$ line cores cannot be interpreted in terms of one-component magnetic field models, suggesting that the magnetic structuring of the solar chromosphere could be substantially more complex than previously thought. For instance, in the presence of a topologically complex distribution of ``weak'' solar magnetic fields, the D$_2$ line core would respond mainly to the strongest and preferentially radially oriented fields, while the D$_1$ line to the weakest and more randomly oriented fields. It remains to be seen whether or not this conclusion is validated after we take fully into account radiative transfer effects and the role of {\\em dichroism} (\\citealt{JTB97}) on the emergent polarization of the ``enigmatic'' Na {\\sc i} D-lines."
    },
    "0201/astro-ph0201206_arXiv.txt": {
        "abstract": "{ We observed \\object{SN\\,1979C} in \\object{M100} on 4 June 1999, about twenty years after explosion, with a very sensitive four-antenna VLBI array at the wavelength of $\\lambda$18\\,cm. The distance to M100 and the expansion velocities are such that the supernova cannot be fully resolved by our Earth-wide array. Model-dependent sizes for the source have been determined and compared with previous results.  We conclude that the supernova shock was initially in free expansion for 6$\\pm2$\\,yrs and then experienced a very strong deceleration. The onset of deceleration took place a few years before the abrupt trend change in the integrated radio flux density curves. We estimate the shocked swept-up mass to be $M_{\\rm sw} \\sim1.6 M_{\\sun}$, assuming a standard density profile for the CSM. Such a swept-up mass for \\object{SN\\,1979C} suggests a mass of the hydrogen-rich envelope ejected at explosion no larger than $M_{\\rm env} \\sim 0.9 M_{\\sun}$. If  \\object{SN\\,1979C} originated in a binary star the low value of $M_{\\rm env}$ suggests that the companion of the progenitor star stripped off most of the hydrogen-rich envelope mass of the presupernova star prior to the explosion. ",
        "introduction": "}  Supernova \\object{SN\\,1979C} in \\object{M100} was discovered on 1979 April 19 by Gus E.\\ Johnson (Mattei \\et\\ \\cite{mat79}) at around magnitude 12 in the visible, and it is thought to have exploded on April 4, 1979 (Weiler \\et\\ \\cite{wei86}). Optical spectra first showed a featureless continuum, which later evolved to exhibit strong H$\\alpha$ emission (e.g., Schlegel \\cite{sch96a}). \\object{SN\\,1979C} was extraordinarily luminous, $M_B^{\\rm max}\\approx -20$ (e.g., Young \\& Branch \\cite{you89}), making it among the most luminous type II supernova ever observed. \\object{SN\\,1979C} was also found to be a radio supernova (RSN)(Weiler and Sramek, 1980). Indeed, at a distance of $16.1\\pm1.3$\\,Mpc estimated for \\object{M100} (Ferrarese \\et\\ \\cite{fer96}) it is one of the intrinsically strongest RSN (Weiler \\et\\ \\cite{wei96}). From H$\\alpha$ widths an expansion speed of 9\\,200$\\pm$500\\,km\\,s$^{-1}$ was estimated (Panagia \\et\\ \\cite{pan80} and G. Vettolani, priv. comm.) for epoch around 45\\,d.  \\object{SN\\,1979C} has been classified as a SN of type II-L. Fesen \\et\\ (1999), hereafter F99, have studied the late-time optical properties of \\object{SN\\,1979C} and other three type II-L supernovae. Late-time optical emission lines from type SNe II-L appear surprisingly steady over long time intervals. These authors find similar expansion velocities in the range 5100--6000\\,km\\,s$^{-1}$ in the strongest lines present (H$\\alpha$, [{\\sc Oi}] 6300,~6364\\,{\\AA}, [{\\sc Oii}] 7319,~7330\\,{\\AA}, and [{\\sc Oiii}] 4959,~5007\\,{\\AA}) for all type II-L supernovae. From the asymmetries of the lines these authors find evidence of dust formation. Strong radio emission from a SN appears to be strongly correlated with late time bright optical emission, with SN\\,1979C as a prime example. In \\object{SN\\,1979C}, all detected lines but H$\\alpha$ show emission peaks at $-5\\,000$ and $-1\\,000$\\,km\\,s$^{-1}$ suggestive of clumpy emission from material coming mostly from ejecta, and physically separated from the material emitting H$\\alpha$. The H$\\alpha$ emission tracks well the peak 6\\,cm emission (Weiler \\et\\ \\cite{wei91}) indicating that both emissions are related to the material from the swept-up shell.  The evolution of the radio emission from SN\\,1979C has been interpreted within the mini-shell model (Chevalier \\cite{che82}) as due to synchrotron emission from the outer part of the shock swept-up shell of circumstellar material attenuated by thermal absorption from the even more distant ionized CSM (Weiler \\et\\ \\cite{wei86}, Montes \\et\\ \\cite{mon00}).  Synchrotron self absorption which has been shown to be relevant in \\object{SN\\,1993J} (Fransson \\& Bj\\\"ornsson \\cite{fra98}, P\\'{e}rez-Torres \\et\\ \\cite{per01}) might also play a role in \\object{SN\\,1979C}, although Chevalier (\\cite{che98}) is not of this opinion. The modulation present in the radio light curves of \\object{SN\\,1979C} has led Weiler \\et\\ (\\cite{wei92}) and Montes \\et\\ (\\cite{mon00}) to model the progenitor as possibly being in a detached eccentric binary system with a less massive companion. >From a study of the emission environment of \\object{SN\\,1979C}, Van Dyk \\et\\ (\\cite{vdy99}) have estimated the mass of the progenitor to be $17-18\\,(\\pm3)$\\,M$_\\odot$. Hydro-dynamical simulations by Schwarz \\& Pringle (\\cite{sch96}) confirm that a binary system similar to that proposed by Weiler \\et\\ (\\cite{wei92}) with a spiral shaped stellar wind around the progenitor is feasible for SN\\,1979C. Such modulation of the density of the gas surrounding the RSN and now being shocked, might have dramatic effects on the shape of the radio structure. A determination of such radio structure would be a powerful way to further understand the pre-supernova phase of the progenitor of \\object{SN\\,1979C}.  \\begin{figure} \\vspace{180pt} \\special{psfile=h3030F1.eps hoffset=-4 voffset=-8 hscale=53 vscale=53} \\caption{Radio light curves for SN\\,1979C in M100 at wavelengths of 20, 6, and 2\\,cm. The curves represent the best-fit model light curves. The best-fit parameters were determined using only data through 1990 December (day $\\sim$4300). Notice the flux density increase from day 4000 onwards (Montes et al. 2000). \\label{fig:lightcurve}} \\end{figure}  The radio emission of this RSN, monitored by Montes \\et\\ (\\cite{mon00}), has started to show an interesting trend which can be seen in Fig.\\ \\ref{fig:lightcurve}: after years of steady decline the radio emission has been constant, or perhaps increasing, since about 1990 ($\\sim10$\\,yr) in a way somewhat reminiscent of the radio emission increase for the case of \\object{SN\\,1987A} (Ball \\et\\ \\cite{ball95}). It appears that after a standard RSN radio emission decline phase (Weiler \\et\\ \\cite{wei96}) the expanding shock may have encountered denser ionized gas which has boosted the radio emission. These results may be in accord with those of \\cite{fes99} which show that high electron densities (10$^{5-7}$ cm$^{-3}$) are necessary to explain the absence of several lines, such as [{\\sc Oii}] 3726,~3729\\,{\\AA} in their optical spectra.  Soft X-rays have been detected from \\object{SN\\,1979C} by Immler \\et\\ (\\cite{imm98}), Ray \\et\\ (\\cite{ray01}), and Kaaret (\\cite{kaa01}). The soft X-ray luminosity appears constant or slightly declining from epoch 16\\,yr through 20\\,yr. This result appears compatible with the model predictions of circumstellar interactions (Chevalier \\& Fransson \\cite{che94}).  VLBI observations at 3.6, 6, 13 and 18\\,cm by Bartel \\et\\ (1985), hereafter B85, determined a size of \\object{SN\\,1979C} and its growth for epochs ranging from Dec~1982 to May~1984 assuming a synchrotron emitting optically thin spherical model. At those early epochs the source could not be resolved by any VLBI array. At an average expansion speed of, for example, $\\sim$8\\,000 km\\,s$^{-1}$, one should expect a shell radius growth rate of $\\sim$ 0.3$\\mu$as\\,day$^{-1}$ which over 20 years should yield a shell radius of $\\sim$2.2\\,mas which could be resolved by $\\lambda$6\\,cm VLBI if detected on the long baselines. However, due to sensitivity limitations of the available arrays, the chances of detecting it at $\\lambda$\\,6\\,cm at the present time are null. On the other hand, at $\\lambda$18\\,cm the emission strength of \\object{SN\\,1979C} and the sensitivity of the arrays make the detection possible. We have carried out state-of-the-art $\\lambda$18\\,cm observations in an attempt to learn as much as possible about the source structure. ",
        "conclusions": "}   Marcaide \\et\\ (\\cite{mar95a,mar95b,mar97}) determined for \\object{SN\\,1993J} a shell structure with a shell width 30\\% of the outer radius, somewhat larger than that predicted by Chevalier \\& Fransson (\\cite{che94}). Therefore, we have adopted a shell of width 30\\% of the external radius as a model for SN\\,1979C in this paper and all our discussion will be based on such model. For the present discussion all previous results by \\cite{bar85} will be transformed (normalized) to this model.  At a distance of 16.1\\,Mpc to \\object{M100} an angular size of 1\\,mas corresponds to a linear size of $2.41\\cdot10^{17}$\\,cm. Hence, the source outer radius (1.80\\,mas) corresponds to $4.33\\cdot10^{17}\\,{\\rm cm}$. Since 20.12\\,yr have elapsed since the explosion, the corresponding average supernova expansion velocity is $\\sim$6\\,800\\,km\\,s$^{-1}$.  Since early expansion speeds were estimated at epoch 45\\,d from the H$\\alpha$ lines as $\\sim$9\\,200\\,km\\,s$^{-1}$ it follows that the expansion has been strongly decelerated by the action of the circumstellar matter. How strongly? It is difficult to make a detailed analysis to answer such a question, since intermediate size estimates are not available. Nonetheless, we now discuss our VLBI results, in conjunction with previously published VLBI and optical spectroscopy data, to show that \\object{SN\\,1979C} has indeed suffered a heavy deceleration in its expansion, starting $6\\pm2$ \\,yr after its explosion.   B85 determined for SN\\,1979C a diameter of $2.16^{+0.44}_{-0.50}$\\,mas 5.17\\,yr after explosion, from 1.67\\,GHz VLBI observations using an optically thin uniform sphere model. This value translates into a diameter of $1.56^{+0.31}_{-0.36}$\\,mas for our discussion  shell model. This size corresponds to an average speed for the first 5.17\\,yr of $\\sim$11\\,500\\,km\\,s$^{-1}$. On the other hand, the average speed of the supernova over the first six weeks after explosion as determined from the widths of H$\\alpha$ lines was $\\sim9\\,200\\pm 500$ \\,km\\,s$^{-1}$ (Panagia \\et\\ \\cite{pan80}, G. Vettolani, priv. comm.), which is about 20\\% lower than the average velocity estimated for the first 5 years from the \\cite{bar85} estimates normalized to our discussion model. These two results are nicely consistent. Indeed, in the mini-shell model of Chevalier \\cite{che82}, the radio emission comes from the forward shock, while the optical, UV, and X-ray emissions come mainly from the reverse shock, which moves slower than the forward shock. Hence the velocity ratio determined from the VLBI and optical observations is roughly what is expected for the radii ratio of the forward and reverse shocks, in a model of 20-30\\% shell width (Chevalier \\& Fransson \\cite{che94}, Marcaide \\et\\ \\cite{mar97}) if \\object{SN\\,1979C} followed a free, self-similar expansion for the first 5 years after its explosion. A free expansion for the first 4 years, followed by a decelerated expansion would also be compatible with the mentioned results. These results should be understood as evidence of free expansion of \\object{SN\\,1979C} for the first 4-5 years and of the presence of a shell-like structure, and not as evidence of a given shell width, since different model normalizations of the results of B85 would yield equally consistent results and a similar conclusion.  An estimate of $m$ ($R\\propto t^m$) using the size estimate of \\cite{bar85} and our size estimate results in $m=0.62^{+0.22}_{-0.17}$. This results shows that \\object{SN\\,1979C} has indeed dramatically decreased its growth rate after a likely initial phase of free expansion ($m \\approx 1.0$; see Fig.\\ \\ref{fig:fit2points}). The optical data, although as scarce as the VLBI data, can be used as a check of the reliability of our estimate of $m$. In a self-similar expansion scenario $R\\propto t^m$ and, consequently, $v\\propto t^{m-1}$. If we now assume that the supernova did not decelerate for the first 5.17 years, we can take the value of 9\\,200 \\,km\\,s$^{-1}$ as the expansion velocity just prior to the phase of strongly decelerated expansion. Combining this value with that of 6\\,200$\\pm$300 \\,km\\,s$^{-1}$ at epoch 14.09 \\,yr (F99 and R. Fesen, priv. comm.), gives $m=0.61^{+0.10}_{-0.11}$, in agreement with our estimate of $m=0.62^{+0.22}_{-0.17}$.   \\begin{figure} \\vspace{174pt} \\special{psfile=h3030F3.eps hoffset=-20 voffset=-25 hscale=65 vscale=65} \\caption{Angular diameter vs.\\ age of \\object{SN\\,1979C}. {\\bf B85} indicates the size estimate from Bartel \\et\\ (1985) normalized to our model (see text), {\\bf F99}, the size estimated by us based on the optical results by Fesen \\et\\ (1999), and {\\bf M02} indicates our results.  The solid line indicates a possible expansion, which is free for the first 5\\,yrs and decelerated from then on. The size estimate {\\bf F99} shown is compatible with this expansion scenario and it would be slightly different for scenarios where the free expansion would last much longer than 5\\,yr, which is not the case. See text. \\label{fig:fit2points}} \\end{figure}    Hence, $m=0.62$ appears to be an adequate value to characterize a phase of strong deceleration in the expansion of SN\\,1979C, assuming it started about 5.17 years after its explosion. If the supernova started to decelerate later than 5.17yr, then $m$ should be smaller, and correspondingly stronger the deceleration of SN\\,1979C. It seems somewhat arbitrary to assume that the deceleration set in at the same time as the observations of B85. Indeed, there is nothing apparent in the integrated radio light curves which could point to a change at that particular epoch. It is intriguing, though, that the integrated radio light curves of SN\\,1979C (Fig. 1) seem to show an increase in flux at around t=2200 days, or 6 yr after its explosion. One could speculate that the flux rise at this epoch is somehow related to a change in the deceleration parameter, likely due to the shock front of the supernova entering a high-density region of the circumstellar medium (CSM).   Our estimate of the angular diameter for SN\\,1979C can also be used to obtain an upper limit for the epoch at which the phase of strong deceleration started. We find that the latest epoch of free expansion which is still compatible with both the optical observations of F99 (epoch 14.09yr) and our VLBI determination at 20.12 yr, and with the fact that a lower limit is set by SN\\,1979C entering a Sedov phase ($m=0.4$), corresponds to $\\sim$8 yr.   The integrated radio emission curves have been well modeled for the first 10 years by  Weiler \\et\\ (\\cite{wei92}) and Montes \\et\\ (\\cite{mon00}) who assumed self-similar expansion and a modulation in the density of the CSM due to the binary interaction, to explain the modulation in the flux density curve. Montes \\et\\ (\\cite{mon00}) found that starting around epoch 10\\,yr the radio emission started to grow after a long decline predicted by the radio emission model. Can our VLBI observations shed any light on what happened at epoch 10\\,yr? The VLBI results by \\cite{bar85} and us require a strong deceleration of the expansion. Such deceleration is compatible with the optical results by F99 for a self-similar expansion. Our results rule out any initial free expansion longer than 8 years. Thus, the onset of decelerated expansion took place well before the epoch 10\\,yr where the radio emission decline turned onto a slight growth. The sooner the onset of the decelerated expansion the milder the deceleration needed to be compatible with the observations. The epoch of that onset is restricted to the range 4-8 years and the corresponding deceleration parameter $m$ to the range 0.67-0.44. (Strictly speaking an onset at 3 years with a corresponding $m$=0.72 is not excluded, but it appears unlikely since it yields a radio to optical velocity ratio of 0.68 at epoch 5.17 yr. Such a ratio is small,  but not yet inconsistent with the models used. In any case, the deceleration required for an early onset is comparable to the strongest decelerations reported in other cases by Marcaide \\et\\ (\\cite{mar97}) for SN\\,1993J in M81 and Mc Donald \\et\\ (\\cite{mcd01}) for 43.31+592 in M82). An onset of decelerated expansion at epoch 6\\,yr (corresponding to a deceleration parameter $m$=0.57) appears most likely and it may be associated with a small halt in the otherwise decreasing flux density trend, a halt which may be partly related to encountering a denser CSM.   If the expansion of SN\\,1979C has significantly decelerated, as our VLBI estimate shows, then the mass of the CSM swept up by the shock front, $M_{\\rm sw}$, must be comparable to or larger than the mass of ejected hydrogen-rich envelope, $M_{\\rm env}$. Our estimated angular diameter translates into a forward shock radius of $\\sim 4.33 \\times 10^{17}$ cm. For a presupernova wind of velocity $v_{\\rm w}$=10\\,km\\,s$^{-1}$, this radius corresponds to $\\sim$13\\,700\\,yr prior to the explosion of \\object{SN\\,1979C}. Weiler et al. (1991) estimated a mass-loss rate of $\\sim1.2 \\times 10^{-4} M_{\\sun}$ yr$^{-1}$ for the progenitor of \\object{SN\\,1979C}. If we now assume that the circumstellar medium up to this distance is a standard one ($\\rho_{\\rm cs} \\propto r^{-2}$), we then obtain $M_{\\rm sw} \\sim 1.6 M_{\\sun}$. Momentum conservation therefore implies that $M_{\\rm env} \\approx 0.59 M_{\\rm sw}$, which results in a mass for the ejected hydrogen-rich envelope of $\\sim 0.94 M_{\\sun}$.  In the case of the Type II supernova SN1993J, whose progenitor was a $\\sim 15 M_{\\sun}$ star in a binary system, mass exchange with its companion meant that the mass of the hydrogen-rich envelope ejected at the explosion was $0.2 - 0.4 M_{\\sun}$ (Woosley et al. 1994, Houck \\& Fransson 1996), much less than the $\\approx 3.3 M_{\\sun}$ one would expect from a single-star model (H\\\"oflich  et al. 1993). Weiler et al. (1986) have suggested that SN1979C was born in a binary system consisting of a 15-18 $M_{\\sun}$ red supergiant and a 10 $M_{\\sun}$ B1 main-sequence star. If we assume that SN~1979C lost, like SN~1993J, most of its hydrogen-rich envelope prior to explosion, then the $\\sim 0.9 M_{\\sun}$ decelerated ejecta suggested from our observations represents at least all of the remaining hydrogen-rich envelope.  >From VLBI observations of SN\\,1979C more than 20 years after its explosion, we conclude that an initial free expansion of \\object{SN\\,1979C} was followed by a strong deceleration which set on $\\sim$6\\,yr after the supernova explosion. We estimate the shocked swept-up mass to be $M_{\\rm sw} \\sim1.6 M_{\\sun}$, assuming a standard density profile for the CSM. Such a swept-up mass for \\object{SN\\,1979C} suggests a mass of the hydrogen-rich envelope ejected at explosion as low as $M_{\\rm env} \\sim0.94 M_{\\sun}$ or lower. If, as suggested by Weiler et al. (\\cite{wei86}) and Schwarz \\& Pringle (\\cite{sch96}) \\object{SN\\,1979C} originated in a binary star (alike \\object{SN\\,1993J}) this value of $M_{\\rm env}$ suggests that the companion of the progenitor star stripped off most of the hydrogen-rich envelope mass of the presupernova star prior to its explosion. New observations at $\\lambda$18\\,cm around 2005 should be able to test the above conclusion and would also be essential to image the true structure of \\object{SN\\,1979C}. On the other hand, a continued rise in the level of flux density at $\\lambda$6\\,cm and improved VLBI instrumentation and techniques like phase-referencing may allow an earlier determination of the intrinsic structure of \\object{SN\\,1979C} with high resolution and a test of the binary star scenario which the present results tend to favor."
    },
    "0201/astro-ph0201083_arXiv.txt": {
        "abstract": "{ A three-dimensional (3D) MHD model is applied to simulate the evolution of a large-scale magnetic field in a barred galaxy possessing a gaseous halo extending to about 2.8~kpc above the galactic plane. As the model input we use a time-dependent velocity field of molecular gas resulting from self-consistent 3D N-body simulations of a galactic disk. We assume that the gaseous halo rotates differentially co-rotating with the disk or decreasing its velocity in the Z direction. The dynamo process included in the model yields the amplification of the magnetic field as well as the formation of field structures high above the galactic disk. The simulated magnetic fields are used to construct the models of a high-frequency (Faraday rotation-free) polarized radio emission that accounts for effects of projection and limited resolution, and is thus suitable for direct comparison with observations.\\\\ We found that the resultant magnetic field correctly reproduces the observed structures of polarization B-vectors, forming coherent patterns well aligned with spiral arms and with the bar.  The process initializing a wave-like behavior of the magnetic field, which efficiently forms magnetic maxima between the spiral arms, is demonstrated. The inclusion of the galactic halo constitutes a step towards a realistic model of galactic magnetic fields that includes as many dynamical components as needed for a realistic description. ",
        "introduction": "Barred galaxies offer a unique opportunity to study the interrelations between gas flows and magnetic field structure.  The bars are known to excite strong spiral patterns in the interstellar gas  and to cause strongly non-axisymmetric gas flows (Englmaier \\& Gerhard \\cite{engl97}). Recent observations show that the magnetic field is strongly modified by gas shearing motions, although the field perturbations do not coincide with the suspected position of gas velocity perturbations, as expected from simple MHD models. Some concepts suggest that the non-axisymmetric gas flows in the bar may constitute the main mechanism amplifying the galaxy-scale magnetic fields (Lesch \\& Chiba \\cite{les97}). Studying strong gaseous spiral patterns driven by the bar perturbation also offer a unique opportunity to understand the relationships between spiral arms and magnetic fields that are currently lively debated.  The question of the gas reaction to the bar potential has a long history (Moss et al. \\cite{mosetal98}, Moss et al. \\cite{mos99} and references therein).  However, little has been done to study the evolution of magnetic field in such objects. Moss (\\cite{mos98}) analyzed the magnetic field evolution in barred galaxies, however his models are restricted to two dimensions, while the vertical structures are essential for the magnetic field evolution. Regular magnetic fields are not only passively deformed by large-scale gas motions in the galactic plane but their evolution is also driven by turbulent motions via the dynamo mechanism. The latter process generates galaxy-scale poloidal (hence also vertical) magnetic field components. Though this genuinely three-dimensional magnetic field structure can still be analyzed in two dimensions by applying a solenoidality condition, we find it of great interest to apply a fully 3D model of the field evolution. Moreover we also get an essential vertical component of the large-scale gas motion, which requires a completely three-dimensional treatment. A three-dimensional analysis of the interaction of the dynamo process with spiral arms has been performed by Rohde \\& Elstner (\\cite{roh98}) and Rohde et al. (\\cite{roh99}).  However, they considered only the turbulence enhancement in star-forming parts of spiral arms with no compression effects included, which is appropriate for the weak density wave objects (NGC 6946, Beck \\& Hoernes~\\cite{bec96}) they intended to explain. A comprehensive study of interactions of magnetic fields with bar-induced gas flows and spiral arms involving the dynamo process (hence carried out in three dimensions) is still urgently needed.  In our previous work (Elstner et al. (\\cite{els00}), v. Linden et al. (\\cite{lin98}), Elstner et al. (\\cite{els98})) we analyzed the magnetic field evolution under the influence of the bar and spiral arms.  However, these studies assumed a thin gaseous disk only. A large gaseous halo may be of crucial importance for the problem.  The dynamo process includes the generation of large-scale poloidal fields then transformed into azimuthal ones by differential rotation.  A realistic analysis of the magnetic field evolution under a combined effect of the gas flows and the turbulent dynamo needs to involve an extended gaseous halo.  For these reasons we performed the study of magnetic field evolution assuming a galaxy composed of a disk and of a rarefied ionized halo with a scale height of 2.8~kpc. This expands our previous results involving a thin disk only. As previously, we use a realistic model of the gas flow, by analyzing motions of inelastically-colliding gas clouds in the self-consistent calculated gravitational potential of stars and gas. We adopted conditions suitable for the formation of a stellar bar. We assumed the magnetic field to be partially coupled to the gas via the turbulent diffusion process. We checked the results for all the assumptions by varying the diffusion coefficient, the strength of the dynamo action, and also by switching off the turbulent dynamo.  The magnetic field in external galaxies is usually studied by means of the radio polarization, yielding the view of the magnetic field integrated along the line of sight and convolved to a certain radio telescope beam. As our models yield complex three-dimensional magnetic field structures, we decided to analyze our magnetic field structures by simulating the polarization maps, which yields much more concise information, easy to compare with observations.  \\noindent ",
        "conclusions": "The evolution of large-scale galactic magnetic fields in a galaxy with a gaseous halo and a bar has been demonstrated using 3D numerical simulations. In order to solve the dynamo equation we used velocities obtained from a self-consistent N-body, sticky particle code. We found that:  \\begin{itemize} \\item[1.]{The magnetic arms  persist in the presence of a large halo enabling a large-scale three-dimensional field.} \\item[2.]{The magnetic arms are also present in interarm regions, in agreement with observations.} \\item[3.]{Our calculations confirm the fact, known from earlier papers, that the pattern speed of the bar is higher than the pattern speed of spiral arms.} \\item[4.]{The magnetic field may take the form of a ``magnetic wave'' with magnetic arms drifting into the interarm areas due to dynamical dissipation and shear amplification independent of the value of the turbulent diffusion.} \\item[5.]{The magnetic pitch angles for all models keep the mean value between $-15\\degr$ and $-25\\degr$ independent of the diffusivity.} \\item[6.]{The edge-on view of the modeled magnetic field agrees with the majority of real observations.} \\end{itemize}"
    },
    "0201/astro-ph0201426_arXiv.txt": {
        "abstract": "We present deep 850~$\\mu$m maps of three massive lensing clusters, A370, A851, and A2390, with well-constrained mass models. Our cluster exposure times are more than 2 to 5 times longer than any other published cluster field observations. We catalog the sources and determine the submillimeter number counts. The counts are best determined in the 0.3 to 2~mJy range where the areas are large enough to provide a significant sample. At 0.3~mJy the cumulative counts are $3.3_{1.3}^{6.3}\\times 10^4$~deg$^{-2}$, where the upper and lower bounds are the 90\\% confidence range. The surface density at these faint count limits enters the realm of significant overlap with other galaxy populations.The corresponding percentage of the extragalactic background light (EBL) residing in this flux range is about $45-65$\\%, depending on the EBL measurement used. Given that $20-30$\\% of the EBL is resolved at flux densities between 2 and 10~mJy, most of the submillimeter EBL is arising in sources above 0.3~mJy. We also performed a noise analysis to obtain an independent estimate of the counts. The upper bounds on the counts determined from the noise analysis closely match the upper limits obtained from the direct counts. The differential counts from this and other surveys can reasonably be described by the parameterization $n(S)=3\\times 10^4$~deg$^{-2}$~mJy$^{-1}/(0.7 + S^{3.0})$ with $S$ in mJy, which also integrates to match the EBL. ",
        "introduction": "\\label{secintro}  The cumulative emission from all sources lying beyond the Galaxy, the extragalactic background light (EBL), provides important constraints on the integrated star formation and accretion histories of the Universe. {\\it COBE} measurements of the EBL at far-infrared (FIR) and submillimeter wavelengths (e.g., \\markcite{puget96}Puget et al.\\ 1996; \\markcite{fixsen98}Fixsen et al.\\ 1998) indicate that the total radiated emission that is absorbed by dust and gas and then reradiated into the FIR/submillimeter is comparable to the total measured optical EBL. Deep submillimeter surveys with SCUBA (\\markcite{holland99}Holland et al.\\ 1999) on the James Clerk Maxwell Telescope\\altaffilmark{5}\\altaffiltext{5}{The JCMT is operated by the Joint Astronomy Centre on behalf of the parent organizations the Particle Physics and Astronomy Research Council in the United Kingdom, the National Research Council of Canada, and The Netherlands Organization for Scientific Research.} have uncovered the brighter obscured sources which give rise to a substantial part of the FIR/submillimeter EBL. The properties of these sources are similar to the properties of the most luminous systems observed locally, the ultraluminous infrared galaxies (ULIGs, \\markcite{sanders96}Sanders \\& Mirabel 1996).  Blank field SCUBA surveys have resolved sources over the $2-10$~mJy range that account for $20-30$\\% of the 850~$\\mu$m EBL (e.g., \\markcite{barger98}Barger et al.\\ 1998; \\markcite{hughes98}Hughes et al.\\ 1998; \\markcite{eales99}Eales et al.\\ 1999, 2001; \\markcite{bcs99}Barger, Cowie, \\& Sanders 1999a; \\markcite{scott01}Scott et al.\\ 2002; \\markcite{borys02}Borys et al.\\ 2002; \\markcite{webb02}Webb et al.\\ 2002). \\markcite{bcs99}Barger et al.\\ (1999a) found that their cumulative source counts per square degree were well described by the power-law parameterization of the differential counts  $$n(S)=N_0/(a+S^{\\alpha})$$  \\noindent with $S$ in mJy, $\\alpha=3.2$, and $N_0=3.0\\times 10^4$~deg$^{-2}$~mJy$^{-1}$. Assuming only that their parameterization provided an appropriately smooth continuation to fluxes below 2~mJy, \\markcite{bcs99}Barger et al.\\ (1999a) constrained their fit to match the EBL with $a\\sim 0.5$. Using this empirical parameterization, they predicted that very approximately 60\\% of the EBL at 850~$\\mu$m should be resolved into discrete sources between 0.5 and 2.0~mJy.  Blank field SCUBA surveys cannot reach the required sensitivities to directly detect the dominant population of $<2$~mJy extragalactic sources due to confusion noise resulting from the coarse resolution of SCUBA (e.g., \\markcite{h01}Hogg 2001). Thus, in order to search for this population with SCUBA, one must observe fields with massive cluster lenses to take advantage of both gravitational amplification by the lens and reduced confusion noise. \\markcite{smail97}Smail, Ivison, \\& Blain (1997) and \\markcite{smail98}Smail et al.\\ (1998) pioneered this method, making the first SCUBA observations of cluster lenses. They studied seven clusters with well-constrained lens models where it is possible to correct the observed source fluxes for lens amplification. Their survey was designed to detect the brightest submillimeter sources in relatively short integration times, and hence it is quite shallow with $3\\sigma$ flux limits around 6~mJy.  Despite the shallowness of the survey, \\markcite{blain99}Blain et al.\\ (1999) tried to determine the number counts at and below 1~mJy, arguing that the small regions of high amplification (around a factor of 6 at 1~mJy and a factor of 24 at 0.25~mJy) could be used to probe these faint fluxes. (A later analysis of another cluster sample of similar depth by \\markcite{chapman02}Chapman et al.\\ (2002) did not extend the counts below 1~mJy.) However, direct inversion at these amplification levels is complicated because redshift uncertainties and small positional changes can have large effects on the amplification. In SCUBA surveys the positions of the submillimeter sources are relatively poorly determined and the redshifts are often unknown. While these effects have comparatively little effect on the shape of the determined counts, they do effectively limit the flux level to which the counts can be considered to be determined. We shall discuss this point further in \\S~4.  The sub-mJy counts are best addressed by obtaining much deeper images, since at lower amplifications there is considerably less redshift and positional sensitivity. At typical amplifications of 1 to 4, the counts can be robustly investigated down to a few tenths of a mJy from images with $3\\sigma$ limits of 1.5 to 2~mJy. Because of this key point, and because the observed areas at the sub-mJy fluxes of interest rise rapidly with deeper obervations, it is a natural and important progression to pursue the faint submillimeter counts with much longer integrations on lensing cluster fields. This is the subject of the present paper. ",
        "conclusions": "\\label{secdisc}  We summarize the current results and compare them with wide-field blank surveys in Fig.~\\ref{fig8}. The present data are shown by the solid squares. We show the lensing analysis of \\markcite{blain99}Blain et al.\\ (1999) as open circles and that of \\markcite{chapman02}Chapman et al.\\ (2002) as open downward pointing triangles. The present analysis shows good overlap with that of Blain et al.\\ at the brighter fluxes, though both are slightly lower than that of Chapman et al. In order to show the counts at bright fluxes we have plotted the blank field surveys of Barger, Cowie, \\& Sanders (1999), Hughes et al.\\ (1998), Eales et al.\\ (2000), Scott et al.\\ (2002), and Borys et al.\\ (2002). The large area survey of Scott et al.\\ and the Hubble Deep Field scan map point of Borys et al.\\ produce slightly shallower counts than the other surveys.  We show a simple broken power-law fit to the counts in Fig.~\\ref{fig8} where the differential counts are given by  \\begin{equation} n(S)=N_0/S^{\\alpha} \\end{equation}  \\noindent with $N_0=2.5\\times 10^{4}$~deg$^{-2}$~mJy$^{-1}$ and $\\alpha=3.0$ above $S=3$~mJy and $N_0=1\\times 10^{4}$~deg$^{-2}$~mJy$^{-1}$ and $\\alpha=2.2$ below 3~mJy. The EBL in this representation diverges at the faint end. We plot the Barger et al.\\ (1999) parametric fit for two values of $\\alpha$ (dashed line shows 3.0, dotted line shows 3.2), $N_0=3\\times 10^{4}$~deg$^{-2}$~mJy$^{-1}$, and $a=0.7$ and $a=0.5$, respectively. Both curves give an integrated surface brightness set to match the EBL measurement of \\markcite{fixsen98}Fixsen et al.\\ (1998) and provide reasonable descriptions of the counts, though the shallower model may be preferred as a better match to the \\markcite{scott02}Scott et al.\\ (2002) data.   \\begin{inlinefigure} \\psfig{figure=cowie.fig8.ps,angle=90,width=3.5in} \\vspace{6pt} \\figurenum{8} \\caption{ A summary of the counts and 90\\% confidence limits between 0.2 and 3~mJy from the analysis of \\S~5 is shown by the solid squares and associated uncertainties. The upper limit at 0.1~mJy, based on the noise analysis of \\S~6, is shown as a downward pointing arrow. The open circles show the points from Blain et al.\\ (1999). These have been slightly displaced to larger fluxes to distinguish them from the present counts. The Blain et al.\\ point at 0.5~mJy is shown as a $2\\sigma$ upper limit for consistency with the present analysis. The open downward pointing triangles show the cluster lensing analysis counts of Chapman et al.\\ (2002). All of the lensing analysis data is only shown below a flux of 5~mJy. We include the wide-field counts of Barger, Cowie, \\& Sanders (1999) (open diamonds), Hughes et al.\\ (1998) (crosses), Eales et al.\\ (2000) (open squares), Scott et al.\\ (2002) (open triangles), and Borys et al.\\ (2002) (asterisk), but we only show $1\\sigma$ uncertainties for the Scott et al.\\ data to avoid confusing the plot; the uncertainties on the other surveys are larger. We show the Barger, Cowie, \\& Sanders (1999) parametric fit (dotted line) and an alternate version (dashed line) of this parametric fit discussed in the text, which provides a better match to the Scott et al.\\ data; it also integrates to match the EBL. We also show a broken power-law representation of the data with slope $-2$ above 3~mJy and $-1.2$ below 3~mJy. This representation is divergent at faint fluxes and must turn over further at some point. \\label{fig8} } \\addtolength{\\baselineskip}{10pt} \\end{inlinefigure}  The contribution to the EBL in the 0.3 to 2~mJy flux range can be directly measured from our counts and is $2.0_{0.8}^{3.2}\\times 10^4$~mJy~deg$^{-2}$, where the upper and lower bounds are the 90\\% confidence range. Thus, the percentage of the EBL residing in this range is $65_{26}^{100}$\\%, if we adopt the 850~$\\mu$m EBL measurement of $3.1\\times 10^4$~mJy~deg$^{-2}$ from \\markcite{puget96}Puget et al.\\ (1996), or $45_{18}^{72}$\\%, if we adopt the measurement of $4.4\\times 10^4$~mJy~deg$^{-2}$ from \\markcite{fixsen98}Fixsen et al.\\ (1998). Given that $20-30$\\% of the EBL is resolved at fluxes between 2 and 10~mJy (\\markcite{bcs99}Barger et al.\\ 1999a; \\markcite{eales99}Eales et al.\\ 1999, 2001, \\markcite{scott02}Scott et al. 2002), it appears that most of the submillimeter EBL is arising in sources brighter than 0.3~mJy."
    },
    "0201/astro-ph0201118_arXiv.txt": {
        "abstract": "A review of the spectroscopic tools needed to characterize AGNs is presented. This review focusses on ultraviolet, optical and infrared emission-line diagnostics specifically designed to help differentiate AGNs from starburst-dominated galaxies. The strengths and weaknesses of these methods are discussed in the context of on-going and future AGN surveys. ",
        "introduction": "The first decade of the 21st Century promises to become the Golden Age of extragalactic astronomy. The 2dF and Sloan Digital Sky Surveys have already made significant contributions to our knowledge of the extragalactic universe. On-going and planned wide-field (pencil-beam) imaging and spectroscopic surveys with 4m (8m)-class telescopes from the ground and in space (e.g., SIRTF, SOFIA, Herschel, NGST) will nicely complement these large-scale surveys and should go a long way to answer some of the most fundamental questions in extragalactic astronomy: How do galaxies form? How do they evolve? How do supermassive black holes fit in this picture of galaxy formation? Which objects are the main contributors to the overall energy budget of the universe? To properly answer these questions, one will need to differentiate objects powered by nuclear fusion in stars (i.e.~normal and starburst galaxies) from objects powered by mass accretion onto supermassive black holes (quasars and AGNs). A wide variety of diagnostic tools have been used in the past for this purpose with different degree of success.  Due to space limitations, the present discussion focusses on {\\em emission-line} diagnostics.  The fundamental principles behind these diagnostics are reviewed in \\S 2. Next, the main diagnostic tools available in the ultraviolet, optical, and infrared domains are described in \\S 3, \\S 4, and \\S 5, respectively. A table listing the main diagnostic lines is given in each of these sections. Additional factors which may complicate the use of these tools are discussed in \\S 6. A summary is given in \\S 7 along with an outlook on the future. Note that this review is not meant to be exhaustive; it is meant to emphasize the practical aspects of starburst/AGN spectral classification.  Readers who are looking for a more detailed discussion of the physics behind these diagnostic tools should refer to the original papers listed in the text. ",
        "conclusions": "UV--Optical--IR emission-line ratios are powerful diagnostics tools to discriminate between starbursts and AGNs. The following ratios have been shown to be the most reliable tools for this purpose. \\begin{itemize} \\item[1.] Ultraviolet: N V $\\lambda$1240/Ly$\\alpha$, N V $\\lambda$1240/He II $\\lambda$1640, C IV $\\lambda$1548/Ly$\\alpha$. \\item[2.] Optical: [O III] $\\lambda$5007//H$\\beta$, [N II] $\\lambda$6583/H$\\alpha$, [S II] $\\lambda\\lambda$6724/H$\\alpha$, [O I] $\\lambda$6300/H$\\alpha$, [O II] $\\lambda\\lambda$7324/H$\\alpha$, [Fe VII] $\\lambda$6087/H$\\alpha$, [Ne V] $\\lambda$3426/H$\\beta$, He II $\\lambda$4686/H$\\beta$. \\item[3.] Near-infrared: Obscured broad Pa$\\alpha$ 1.875 $\\mu$m, [Si VI] 1.962 $\\mu$m/Pa$\\alpha$. \\item[4.] Mid-Infrared: [Ne V] 14 $\\mu$m/[Ne II] 12.8 $\\mu$m, [O IV] 26 $\\mu$m/[Ne II] 12.8 $\\mu$m, EW(PAH 7.7 $\\mu$m), overall SEDs especially 25 $\\mu$m/60 $\\mu$m colors. \\end{itemize} A number of issues complicate the use of line ratios as discriminants between starburst and active galaxies, but additional measures can be used to clarify the situation: \\begin{itemize} \\item[1.] Shock ionization: If shocks are important, one would generally expect correlations between the line ratios and gas kinematics, a UV continuum extended on the same scale as the shock structure, and high gas temperatures. \\item[2.] Aperture effects: One should use a constant linear aperture to avoid variations in the contributions from circumnuclear starbursts. \\item[3.] Morphological bias: The spectral classification is likely to depend on the morphology of the host, especially the merger phase. Selection methods based on morphology will bias the sample. \\item[4.] Metallicity: Massive host galaxies in the local universe have larger metallicity, but high-redshift galaxies should be less dusty and less metal rich. One needs to use emission-line diagnostics which are properly calibrated as a function of metallicity and reddening. \\end{itemize}  Several new instruments will help refine the diagnostic tools discussed in this paper. The Cosmic Origins Spectrograph (COS), to be installed in 2003 on HST, will provide the high ultraviolet throughput needed to calibrate the UV diagnostic tools as a function of metallicity, evaluate the importance of shock ionization with the use of the C III and NIII temperature-sensitive line ratios, and to help resolve the circumnuclear starbursts and shock-excited winds around AGNs.  The advent of SIRTF will help in the calibration of the infrared diagnostic tools as a function of metallicity {\\em and dust extinction}. This spacecraft will also be a powerful instrument to search for infrared-bright AGNs. Ground-based work with adaptive optics and integral-field units will improve the sensitivity of searches for obscured AGNs by focussing on the inner regions of galaxies and avoiding the circumnuclear material associated with other phenomena. Spectroscopic follow-ups from the ground will help identify and classify AGN candidates in space-based and submm-selected samples."
    },
    "0201/astro-ph0201268_arXiv.txt": {
        "abstract": "Spectral absorption features in active galactic nuclei (AGNs) have traditionally been attributed to outflowing photoionized gas located at a distance of order a parsec from the central continuum source. However, recent observations of QSO FIRST J104459.6+365605 by de Kool and coworkers, when intepreted in the context of a single-phase gas model, imply that the absorption occurs much farther ($\\approx 700\\ {\\rm pc}$) from the center. We reinterpret these observations in terms of a shielded, multiphase gas, which we represent as a continuous low-density wind with embedded high-density clouds.  Our model satisfies all the observational constraints with an absorbing gas that extends only out to $\\sim 4\\ {\\rm pc}$ from the central source. The different density components in this model coexist in the same region of space and have similar velocities, which makes it possible to account for the detection in this source of absorption features that correspond to different ionization parameters but have a similar velocity structure. This model also implies that only a small fraction of the gas along the line of sight to the center is outflowing at the observed speeds and that the clouds are dusty whereas the uniform gas component is dust free. We suggest that a similar picture may apply to other sources and discuss additional possible clues to the existence of multiphase outflows in AGNs. ",
        "introduction": "\\label{intro} Modeling active galactic nuclei (AGNs) is a challenging endeavor in part because of the uncertain geometry of the gas and dust within their cores.  Our best understanding of the gas distribution comes from reverberation mappings \\citep{BM82,NP97}, which indicate that the broad emission-line region (BELR) lies at a distance $R_{\\rm BELR} \\approx 0.01 L^{1/2}_{44}\\ {\\rm pc}$ from the central continuum, where $L_{44}$ is the luminosity in the spectral band $\\sim 0.1-1$ {\\micron} in units of $10^{44}$ ergs \\persec.  Based on observed absorption of BELR features by broad absorption-line region (BALR) gas and on optical polarization measurements, it has been inferred that the BALR lies outside of the BELR; the best estimates have placed the BALR somewhere between a few tenths and several parsecs from the central source \\citep[e.g.,][]{T88}.  Velocities for these outflows are well determined from Doppler shifts: the BELR and BALR components have speeds $\\lesssim 5000~{\\rm to}~8000~{\\rm km~s}^{-1}$ and $\\lesssim 30000~{\\rm km~s}^{-1}$, respectively.  Other outflow components have also been detected. Observations of the ``warm absorber'' (a partially ionized X-ray absorbing gas) in Seyfert galaxies and radio-quiet QSOs, and of a UV-absorbing component apparently associated with the warm absorber, indicate gas outflowing at $\\lesssim 10^3~{\\rm km~s}^{-1}$. Furthermore, spectral lines classified as being from the Narrow Line Region (NLR) correspond to velocities of $\\lesssim$ several hundred km~s$^{-1}$.  Different theories have been proposed to explain some of these outflows.  For example, \\citet*{EBS92} considered a hydromagnetically-driven disk outflow of discrete, long-lived clouds as a model of the BELR, whereas \\citet{M95} and \\citet{P00} explored radiatively-driven, continuous disk winds as the origin of the BELR and the BALR. These models all involved a single-phase gas medium.  Recently, \\citet[hereafter dK01]{dK01} analyzed a Keck HIRES spectrum of the FIRST Bright Quasar Source J104459.6+365605 (hereafter, FBQS 1044).  They inferred that excited Fe~II levels are not populated according to local thermodynamic equilibrium. This observation implies an electron number density $n_{e} \\approx 4 \\times 10^{3}~{\\rm cm}^{-3}$, which, together with the observed broad absorption in Mg~II [qualifying this object as a broad absorption-line quasar (BALQSO)], the Mg~I absorption features, and single gas-phase photoionization models, place this BALR at $\\sim$ 700 parsecs from the central source.  Not only does this surprising result contradict the picture of a close-in BALR, but, in addition, dK01's thorough analysis suggests that the gas only partially covers the continuum source.  It is not easy to understand how the absorbing gas could be approximately 1 kpc from the center and yet only partially cover the AGN core. Additionally, the Keck spectra indicate distinct groupings of gas in velocity space; it is also difficult to explain how these apparent clumps could move to such a large distance and still remain bunched up, retaining distinct identities.  Even more striking is dK01's observation that different ionization states have a similar kinematic structure (i.e., residual intensity as a function of velocity): in particular, both Fe~II and Mg~I exhibit absorption troughs at -200, -1250, -3550, and -3800 km~s$^{-1}$.  The $700~{\\rm pc}$ distance estimate results from combining the ionization parameter\\footnote{The ionization parameter $U \\equiv Q/4{\\pi}nr^2c$ is the dimensionless ratio of the hydrogen-ionizing photon density to the hydrogen number density $n$: $Q$ is the rate of incident hydrogen-ionizing photons, $r$ is the distance from the continuum source, and $c$ is the speed of light.} implied by the observed low-ionization Mg and Fe absorption lines with the inferred values of $n_e$ and $L$ (=~$10^{46} {\\rm ergs~s^{-1}}$).  One can reduce this distance estimate by attenuating the incident continuum by an intermediate gas ``shield.'' In this model, however, ions at different ionization states (such as the Fe~II and Mg~I components seen in FBQS 1044) are expected to occupy different regions in space, with lower values of $U$ corresponding to larger distances from the center \\citep[e.g.,][]{VWK93}. This stratification will likely lead to disparate velocities for the different ions: for example, models of self-similar MHD winds and of radiatively-driven constant-$U$ outflows predict that the terminal velocity, $v_\\infty$, varies as $r_{\\rm inject}^{-1/2}$, where $r_{\\rm inject}$ is the radius at which the gas is injected into the outflow \\citep*[e.g.,][]{BP82,A94}. Models that rely solely on shielding and distance to separate the different ionization components therefore cannot explain dK01's observations, in which different ionization states are found to have similar velocities.  We propose to explain the observations of FBQS 1044 by generalizing the single-phase shielded-gas model, attributing the different ionization states to different density components in a multiphase outflow. These components coexist at the same distance from the center and thus have different ionization parameters but essentially the same velocities. In particular, if the high- and low-ionization lines arise, respectively, in a continuous wind and in dense clouds that are embedded in the outflow, then all the absorption components produced in a given region of the wind will exhibit similar kinematic signatures.  In our model, the continuous gas component extends from near the black hole's event horizon out to the distance where the observed Fe~II absorption and electron density can be reproduced. The inner part of this component (interior to the BALR) is identified as the ``shield.'' This region could be associated with an MHD-driven \\citep{KK94}, a Thomson scattering-driven \\citep[e.g.,][]{B01}, or a ``failed'' line-driven \\citep[e.g.,][]{M95,P00} disk wind, or with a disk corona \\citep[e.g.,][]{EBS92}. The outer region of the continuous gas component is outflowing and accounts for the Fe~II and Mg~II absorption, but is still too highly ionized to contain Mg~I absorbing gas. Within this outflow are embedded higher-density clouds that account for the lower-ionization Mg~I absorption: their high density yields a lower U in the clouds, allowing Mg~I to exist.  Such a two-component outflow may arise naturally in the context of a centrifugally driven disk wind, which could uplift clouds from the disk surface by its ram pressure and confine them by its internal magnetic pressure \\citep[e.g.,][]{EBS92,KKE99,EKK01}.  Alternatively, the clouds may represent transient density enhancements that are produced by turbulence \\citep[e.g.,][]{BF01} or by shocks in a radiatively driven wind \\citep[e.g.,][]{A94}.  For illustration, we adopt here the ``clouds uplifted and confined by an MHD wind'' picture. ",
        "conclusions": "The most robust result of our study is that a shielded, multiphase absorption region reproduces the observations of FBQS 1044 on a conventional BALR scale ($\\approx 4$ pc).  In addition, when one attributes the Fe~II and Mg~II absorption to a low-density outflow component and the Mg~I absorption to a cospatial high-density outflow component, it is possible to explain the similar kinematic structure of the respective spectral features.  We also find that only a small fraction of the gas along the line of sight can be outflowing at the observed speeds and that only the high-density component of the outflow is dusty.  We derived these results using a ``clouds embedded in a continuous MHD disk wind'' model, but our conclusions also apply to other plausible scenarios that include a continuum-shielding gas column and an absorption region that contains distinct low- and high-density components. Our basic conclusions appear to be quite general, although the precise composition of the absorbing gas and its detailed spatial and kinematic properties are not fully constrained by the observations and remain model dependent.  In addition to explaining the FBQS 1044 observations, this picture may be relevant to the interpretation of absorption features in similar objects where single-phase models imply a large distance.  For instance, in the case of the radio-loud galaxy 3C 191, absorber distances of $\\sim 28$ kpc were inferred by \\citet{H01} using similar arguments to those employed by dK01.  A multiple-phase model could place these absorbers much closer to the central source.  This interpretation may also be applicable to other AGN observations. As outlined in \\S~1, several distinct outflow components have been inferred in various types of AGNs. There is now growing evidence that these components may not be single-phase. For the warm-absorber component, which has been inferred to give rise to both X-ray and UV absorption \\citep[e.g.,][]{C97,MWE98,MMWE01}, there are indications in at least some sources that the X-ray and UV absorbing components are not identical (e.g., the Seyfert 1 galaxy NGC 3783 --- \\citealt*{K00, KCG01}).  Furthermore, it appears that the UV and X-ray absorbing components are themselves divided into multiple zones. In the UV regime, this has been established in Seyfert 1 galaxies like NGC 3783 \\citep{KCG01} and NGC 3227 \\citep{CKBR01}.  In the X-ray regime, this is exemplified by the Seyfert~1 galaxy MCG-6-30-15, which was modeled by \\citet*{MFR00} as an extended multi-zone medium. These authors suggested that, in reality, these zones may correspond to a continuum of clouds at different radii and different densities, as in the BELR model of \\citet{B95}. A hierarchical, turbulent-gas realization of this concept was recently presented by \\citet{BF01}. Our ``clouds in a continuous wind'' scenario is designed to explicitly address the outflowing nature of the absorbing gas but is otherwise similar to the above picture in its description of spatially coexisting multiple phases. For simplicity, we have treated the clouds as long-lived, pressure-confined entities, but it is entirely conceivable that the clouds are transient features that arise in a turbulent outflow. The same basic picture of a multiphase medium may thus apply to the gas in the BELR, the BALR \\citep[e.g.,][]{A99}, the warm absorber, and even the NLR \\citep[e.g.,][]{K01}.  In conclusion, we have argued that the observations of FBQS 1044 can be interpreted in the context of the standard BALR picture in terms of a gas outflow that consists of (at least) two phases, which we modeled as a continuous wind with embedded dense clouds, shielded from the central continuum.  If the shield is identified with the continuous component, then only a fraction of it can be outflowing at a speed that approaches (or exceeds) the value in the Fe/Mg absorption region. We also deduced that the clouds are dusty but that the shield is effectively dust free.  As far as we are aware, this is the first instance of a BALR outflow in which the data provide direct evidence for the existence of a multiphase medium. Together with other pieces of evidence, this result lends support to the view that all the major outflow components in AGNs may contain multiple phases."
    },
    "0201/astro-ph0201097_arXiv.txt": {
        "abstract": "MACHe3 (MAtrix of Cells of superfluid \\hetro) is a project of a new detector for direct Dark Matter search, using superfluid \\hetrois as a sensitive medium. This paper presents a phenomenological study done with the DarkSUSY code, in order to investigate the discovery potential of this project of detector, as well as its complementarity with existing and planned devices. ",
        "introduction": "A substantial body of astrophysical evidence supports the existence of non-baryonic Dark Matter (DM) in the halo of our galaxy, in particular in the form of new, yet undiscovered, weakly interactive massive particles (WIMPs)~\\cite{jungmanphysrep}. One of the leading candidates is the neutralino predicted by the supersymmetric extensions of the Standard Model of particle physics.\\\\ Following early experimental works~\\cite{lanc}, a superfluid \\hetrois detector has recently been proposed \\cite{firstmac3} for direct Dark Matter search. Monte Carlo simulations have shown that a high granularity detector, a matrix of superfluid \\hetrois cells, would allow to reach a high rejection factor against background events, leading to a low false event rate. The purpose of this paper is to present a full estimation of the neutralino (\\neutt) event rate, within the framework of the Minimal Supersymmetric Standard Model (MSSM), in order to compare with background event rate obtained by Monte Carlo simulation~\\cite{firstmac3,daniel}. Finally, the complementarity with existing devices, both for direct and indirect detection, will be shown.  \\subsection{Experimental device} The elementary component of MACHe3 is the superfluid \\hetrois cell~\\cite{prl95,prb98}. It is a small copper cubic box ($\\mathrm{V} \\simeq 125 \\,{\\rm mm}^3$) filled with superfluid \\hetrois in the B-phase. This ultra low temperature device (T $\\simeq \\!100 \\,\\mu\\mathrm{K}$) presents a low detection threshold ($\\mathrm{E}_{th}\\simeq 1 \\,\\mathrm{keV}$). An experimental test of such a prototype cell has been done at CRTBT in June 2001. Preliminary results \\cite{thesefmayet,santosgamma} show that a threshold value down to $\\sim 1 \\,{\\rm keV}$ has been achieved, and that a stability of the order of one week, at T $\\simeq \\!100 \\,\\mu\\mathrm{K}$, has been obtained.\\\\ The final version of the detector will be a matrix of 1000 cells of $125 \\,{\\rm cm}^3$ each. The idea is to take advantage both on the energy loss measurement and the correlation among the cells to discriminate neutralino events from those of background (neutrons, \\gams and muons). The design of the matrix has been optimized, with a Monte Carlo simulation \\cite{firstmac3}. In the preferred configuration, a $10 \\,{\\rm kg}$ detector, the false event rate has been shown to be as small as $\\sim\\! 10^{-1}\\, {\\rm day}^{-1}$ for neutron events\\footnote{Fast neutron contribution, from interaction of muons in the rock, is expected to be negligible as the enregy release in the cell will be much greater than $6\\,{\\rm kev}$.} and $\\sim\\! 10^{-2}\\, {\\rm day}^{-1}$ for muon events \\cite{firstmac3}. Background from \\gam events needs to be taken into account. Energy loss measurement and correlation among the cells within a $10 \\,{\\rm kg}$ detector allows to obtain a rejection up to 99.8\\% for $\\sim 2\\,{\\rm MeV}$ $\\gamma$-rays \\cite{firstmac3}. Additionnal internal tag on $\\gamma$-rays may be obtained using a new matrix configuration in which two neigbouring cells share a common copper wall, thus greatly improving correlation factor while reducing the amount of copper used. Monte Carlo studies are under way \\cite{santosgamma}. Furthermore, $\\gamma$-ray contamination is to be estimated in the forthcoming months for a multicellular prototype matrix.\\\\ In order to compare these false event rates with the expected neutralino rate, only muons and neutrons will be taken into account in the following. We shall recall that neutrons are usually considered as the ultimate background noise for this type of search, as they interact {\\it a priori} like WIMPs.  \\subsection{\\hetrois as a sensitive medium for direct DM search} Several properties of \\hetrois make this nucleus a promising candidate for a sensitive medium for direct DM search.  {\\bf a)} Concerning background rejection, and as outlined in \\cite{firstmac3,york}, the neutron capture process offers the possibility to discriminate neutron and \\neut event, when considering a ${\\rm 10\\, kg}$ granular detector. Compton cross-section being small ($\\sigma \\lesssim 1\\,{\\rm barn}$), the interactions with \\gams will be minimized. Eventually, as explained in \\cite{prb98}, superfluid \\hetrois is produced with an extremely high purity, the only solute being \\hequatre in a negligible fraction. Consequently, no contamination  from radioactive materials is expected in the sensitive medium. Of course, natural radioactivity from external materials (${\\rm Cu}$, ...) has to be taken into account, by a careful selection of these materials.  {\\bf b)} Concerning neutralino detection the advantage is twofold. First, the maximum recoil energy does only slightly depend on the neutralino mass, due to the fact that the target nucleus (${\\rm m=2.81 \\, GeV}\\!/c^2$) is much lighter than the incoming  \\neut (${\\rm M_\\chi \\geq 32 \\,GeV}\\!/c^2$), considering latest results from collider experiments \\cite{pdg2000}. As a matter of fact, the recoil energy range needs to be studied only below ${\\rm 6 \\,keV}$, see \\cite{firstmac3,thesefmayet}. Second, \\hetrois being a 1/2 spin nucleus, an \\hetrois detector will be sensitive mainly to axial interaction, making this device complementary to existing ones, as shown below. In fact, the axial interaction is largely dominant (up to three orders of magnitude) in all the SUSY region associated with a substantial elastic cross-section \\cite{thesefmayet}.   \\subsection{Theoretical framework} \\label{sec:theo} This phenomenological study has been done with the DarkSUSY code\\footnote{The version used is 3.14.01, with correction of some minor bugs.} \\cite{ds}, within the framework of the phenomenological Supersymmetric Model, namely with the following free parameters : \\begin{equation} \\mu ,\\; \\mathrm{M}_2 ,\\; \\tan \\beta ,\\; m_{A} ,\\; m_{0} \\;\\mathrm{and} \\; A_{b,t} \\end{equation}  \\noindent with $\\mu$ ($\\mathrm{M}_2$) the Higgsino (Gaugino)  mass parameter, $\\tan \\beta$ the ratio of Higgs vacuum expectation values, $m_{A}$ the CP-odd Higgs boson mass, $m_{0}$ the common scalar mass and $A_{b,t}$ the soft trilinear coupling parameters\\footnote{For a good introduction to MSSM models, we refer the reader to \\cite{djouadi}.}.\\\\ Apart from $A_{b,t}$ chosen at fixed zero value, as their influence is expected to be negligible,  all the parameters have been scanned on a large range, with a variable number of steps (tab.~\\ref{tab:scan}). This scan of the free parameters corresponds to a total number of supersymmetric (SUSY) models of the order of $2 \\times 10^6$.\\\\ We shall suppose all through this work the neutralino (\\neutt), the lightest supersymmetric particle, as the particle making up the bulk of galactic cold DM. Each SUSY model is then checked not to be excluded by collider experiments~\\cite{pdg2000}, including $b \\to s\\gamma$ limit.  The nest step, the evaluation of the relic density, is the key point of any Dark Matter calculations. Given the number of free parameters (5 in our case), the allowed SUSY parameter space may be extremely large, leading to a \\neut relic density ranging on up to five orders of magnitude \\cite{thesefmayet}. In order to exclude SUSY models giving a \\neut relic density too far away from  the estimated matter density in the Universe \\cite{omegam} ($\\Omega_{\\mathrm{M}} \\simeq 0.3$), only models with $\\Omega_\\chi$ in the following range are considered : \\begin{equation} 0.025 \\leq \\Omega_\\chi \\mathrm{h_0}^2 \\leq 1 \\label{eq:omega} \\end{equation} \\noindent where $\\mathrm{h_0} = (0.71 \\pm 0.07)\\times^{1.15}_{0.95}$ is the normalized Hubble expansion rate~\\cite{pdg2000}.\\\\ The lower limit comes from the condition that the neutralino relic density has to be at least greater that the baryonic density, and the upper limit is a  conservative limit so that \\neut do not give a density greater than the Universe\\footnote{It can be noticed that selecting on $\\Omega_\\chi \\mathrm{h_0}^2$ allows for a slightly looser selection.}. The SUSY model is thus checked to really provide a good non-baryonic Dark Matter candidate. We follow \\cite{pbarberg} in the choice of the \"cosmologically interesting\" range of $\\Omega_\\chi$. It should be noticed that this loose selection allows to exclude a large number of models in our SUSY scan. For further details concerning the calculation of the \\neut relic density, we refer the reader to \\cite{ds}.  As the detection on earth is concerned, a galactic halo model has to be considered. In the following, standard parameters have been used, in a spherical isothermal halo distribution, with a local density ($\\rho_0$) and an average velocity~($v_0$), with the following values : \\begin{equation} \\rho_0  \\!=\\!   0.3\\,{\\rm GeV}\\!/c^2\\,{\\rm cm}^{-3}\\, \\;\\mathrm{and}\\;  \\,v_0 \\!=\\! 220\\,{\\rm km}\\,{\\rm s}^{-1} \\label{eq:theo} \\end{equation}  \\noindent These parameters are widely used for dark matter detection computations, see \\cite{rick} for instance. No clumpy galactic dark matter structures \\cite{clumps} are considered hereafter, since the effect is expected to be small both for direct detection and neutrino telescopes, in which the signal depends on the local halo density, as emphasized in \\cite{clumps2}. \\vspace*{-7mm} ",
        "conclusions": "It has been shown that a ${\\rm 10\\,kg}$ high granularity \\hetrois detector (MACHe3) would allow to obtain, in many SUSY models, a \\neut event rate higher than the estimated (neutrons and muons) background. MACHe3 would thus potentially allow to reach a large part of the SUSY region, not excluded by current collider limits and for which the neutralino relic density lies within the range of interest. Furthermore, it has been shown that this project of new detector would be sensitive to SUSY regions not covered by future or ongoing DM search detectors, both for direct and indirect detection, thus highlighting the complementarity of MACHe3 with existing or planned devices.     \\noindent \\textbf{Acknowledgments : }\\\\ The authors wish to thank R.~Gaitskell and V.~Mandic for the convenient Dark Matter tools \\cite{rick}. One of us (FM) is grateful to G.~Coignet, Y.~Giraud-H\\'eraud, L.~Mosca, G.~Sajot and C.~Tao, for fruitful discussions on this subject."
    },
    "0201/hep-ph0201107_arXiv.txt": {
        "abstract": "A review of Big Bang Nucleosynthesis (BBN) is presented. Observations of deuterium and helium-4 are discussed. Some BBN restrictions on non-standard physics, especially on neutrino properties and time-variation of fundamental constants are given. \\vspace{1pc} ",
        "introduction": "}  Big Bang Nucleosynthesis (BBN) is known to be one of three solid pillars on which the Standard Cosmological Model (SCM) stands. The other two include General Relativity (GR) and Cosmic Microwave Background Radiation (CMBR). An agreement of BBN calculations of light element abundances with observations presents the strongest proof in favor of the statement that 12-14 billion years ago the universe was indeed hot with the temperatures in MeV range and that the entropy per baryon is huge, about $10^{9}$.  According to the theory, light elements $^2H$, $^3He$, $^4He$, and $^7Li$ have been created in the early universe during first few hundred seconds of her existence. The abundances of these elements span 9 orders in magnitude and are in excellent, good, or reasonable agreement with the observational data, depending upon the moment when the comparison of theory with the data is taken and upon the personal point of view of a researcher.  The theory of BBN is robust, well defined, and quite precise. The largest uncertainty is introduced by the values of the cross-sections of nuclear reactions. Theoretical accuracy is better than 0.1\\% for $^4He$, better than 10\\% for $^2H$ and is about 20-30\\% for $^7Li$~\\cite{fiorentini98,burles99}. In all the cases theoretical uncertainty is much smaller than observational precision. Observations of light elements encounter two serious problems: systematic errors and evolutionary effects. We will discuss them below.  In the next section physics of BBN and essential parameters and inputs are described. In section 3 observational data are analyzed (looking from outside by a non-expert). In section 4 modifications of the standard scenario are discussed. Conclusion is presented in the last section 5. ",
        "conclusions": "We see that gross features of BBN well agree with observations but the latter are not yet sufficiently accurate to make it really a precise science. Moreover there is a trend to discrepancy between the observations of deuterium which indicate a higher value of $\\eta_{10}$ than the observations of $^4 He$. Hopefully it will be clear in a few coming years if this is a real problem or an artifact of systematic and evolutionary effects. We have not discussed above primordial $^7 Li$ because the accuracy of its measurements are rather low now but potentially this element could be very important for verification of the standard model. A recent discussion of $^7 Li$ can be found in ref.~\\cite{olive00}. Still even with the existing level of accuracy BBN permit to put powerful constraints on deviations from the standard model. The number of extra neutrino species allowed by the contemporary observations is about unity and with this bound very little can said about mixing parameters between active and sterile neutrinos. However, if $\\Delta N$ could be reduced, say, to 0.3 the limit would be meaningful. The restriction on the time $\\alpha$ is quite strong and may exclude some of the models predicting such variation."
    },
    "0201/astro-ph0201024_arXiv.txt": {
        "abstract": "We have used available intermediate degree {\\it p}-mode frequencies for the solar cycle 23 to check the validity of previously derived empirical relations for frequency shifts (Jain {\\it et al.}: 2000, {\\it Solar Phys.} {\\bf 192}, 487). We find that the calculated and observed frequency shifts during the rising phase of the cycle 23 are in good agreement. The observed frequency shift from minimum to maximum of this cycle as calculated from MDI frequency data sets is 251 $\\pm$ 7 nHz and from GONG data is 238 $\\pm$ 11 nHz. These values are in close agreement with the empirically predicted value of 271 $\\pm$ 22 nHz. ",
        "introduction": "An understanding of the physical processes responsible for the changes in the solar {\\it p}-mode frequencies could provide an important clue to the inner workings of the solar activity cycle. It is now well established that these frequencies change with time (Woodard and Noyes, 1985) and show a positive correlation with activity indices (Woodard {\\it et al.}, 1991; Bachmann and Brown, 1993). Over the last two decades, attempts have been made to precisely measure  the  changes in {\\it p}-mode oscillation frequencies. More recently, with the Global Oscillation Network Group (GONG) (Harvey {\\it et al.}, 1996)  and Michelson Doppler Imager (MDI)   on board Solar and Heliospheric Observatory (Scherrer {\\it et al.}, 1995)   instruments, the measurements are made   consistently with an accuracy of one part in 10$^5$ or better. These continuous   data sets further confirm that the oscillation frequencies are well correlated with  the activity indices (Jain, Tripathy, and Bhatnagar, 2000 and {\\it references therein}).  Following a different approach and using a data set of eight years  between 1981 and 1989, Rhodes {\\it et al.} (1993) reported that the frequency shifts are also correlated with the change in various activity  indicators. This study was  extended to the rising part of the cycle 23   by Jain {\\it et al.} (2000; hereafter JTBK).  Using the GONG frequencies for the period May 1995 to October 1998, they  confirmed that the frequency shifts are better correlated with the change in  activity  indices. In an attempt to quantify the changes in mode frequencies, JTBK    derived  empirical relations between the shift  in frequencies and change in the level of activity indices and  showed  that these relations do not change significantly from cycle to cycle.   Using a limited data set  from GONG network for the ascending part of the current cycle 23,  it was found that the  calculated and observed frequency shifts were in close agreement.    The motivation of this paper is to check the validity of the derived relations for the maximum and descending phase of the  current solar cycle. This is accomplished by determining the frequency shifts using the smoothed sunspot number and equation (8) of JTBK.  The calculated shifts are compared with the measured shifts obtained from  recent observations. We find that the calculated and observed frequency shifts are  in close agreement confirming  that the derived relations can be reliably used to estimate the {\\it p}-mode frequencies for past, present and future solar activity cycles, if the solar acitivity index is  known. ",
        "conclusions": "We estimate the change in {\\it p}-mode frequencies for solar cycle 23 using the equation (8) of JTBK:   \\begin{eqnarray} \\delta\\nu & =& (2.41 \\pm 0.19)~\\delta R_s - (0.48 \\pm 1.68), \\end{eqnarray} where $\\delta \\nu$ is given in nHz and  $\\delta R_s$ is the change in smoothed sunspot number. The variation of sunspot number taken from the {\\it Solar Geophysical Data} web page for the observing period of MDI is shown in Figure 1. The estimated frequency shifts obtained from Equation (1) are  plotted in Figure 2 for the period from  January, 1995 to December, 2000. The measured shifts for both GONG (lower panel) and MDI (upper panel) data sets are also plotted in the same figure and  shows that the observed  frequency shifts  are in  close agreement with those obtained from the empirical relation. During the ascending phase, the GONG frequency shifts agree better with the derived shifts while the frequency shifts of MDI data sets are in close agreement during the maximum phase of the cycle. On an average,  we find that the deviation between the calculated and observed frequency shifts is within 1$\\sigma$  error level.    \\begin{figure} \\begin{center} \\leavevmode \\input epsf \\epsfxsize=3.75in \\epsfbox{fig3.ps}\\\\ \\leavevmode \\caption{The estimated frequency shifts for the solar cycle 23 using predicted sunspot numbers. The dashed lines represent the  errors in the calculation due to the errors in predicted activity index. } \\end{center} \\end{figure}  Quantitatively, the derived shift between the minimum and maximum of the current solar cycle amounts to 271 $\\pm$ 22 nHz corresponding to the  maximum smoothed sunspot number of  120.8. This can be compared to the observed  shifts of 251 $\\pm$ 7 nHz for MDI and 238 $\\pm$ 11 nHz for GONG data and clearly shows a discrepancy near the maximum phase of the cycle indicating to the complex nature of the relationship that may exist  between activity index and the frequecy shift. Earlier, JTBK had quoted a maximum shift of 265 $\\pm$ 90 nHz corresponding to the predicted maximum sunspot number of 118 $\\pm$ 35.  The estimated frequency shifts for the complete solar cycle 23 (1996 - 2008) are plotted in Figure~3. The solid line represents the estimated shifts calculated using the predicted smoothed sunspot number taken from the {\\it Solar Geophysical Data} web page and dashed lines are the predicted 1$\\sigma$ error.  In summary, we have used available intermediate degree {\\it p}-mode frequencies from GONG and MDI projects for the solar cycle 23 to check the validity of previously derived empirical relations for frequency shifts (Jain {\\it et al.}, 2000). We find that the calculated frequency shifts  are in close agreement with the observed shifts during the period considered in this analysis which includes the rising phase of the cycle 23. We conclude that the empirical relations as derivd by JTBK can be considered as good predictors of frequency shifts for solar activity cycles."
    },
    "0201/astro-ph0201162_arXiv.txt": {
        "abstract": "We have detected a strong periodicity of 1.80$\\pm$0.05 hours in photometric observations of the brown dwarf Kelu-1. The peak-to-peak amplitude of the variation is $\\sim$1.1\\% (11.9$\\pm$0.8\\,mmag) in a 41nm wide filter centred on 857nm and including the dust/temperature sensitive TiO \\& CrH bands. We have identified two plausible causes of variability: surface features rotating into- and out-of-view and so modulating the light curve at the rotation period; or, elliposidal variability caused by an orbiting companion. In the first scenario, we combine the observed $v \\sin i$ of Kelu-1 and standard model radius to determine that the axis of rotation is inclined at 65$\\pm$12\\deg\\ to the line of sight. ",
        "introduction": "The study of rotation and variability in main sequence stars has led to a great improvement in our understanding of their physics (e.g. Stauffer \\& Hartmann 1986). Recently, several groups have shown that variability can also be detected in substellar brown dwarfs (Tinney \\& Tolley 1999; Bailer-Jones \\& Mundt 1999, 2001; Mart\\'\\i n, Zapatero Osorio \\& Lehto 2001).  In this paper we present differential photometry of the brown dwarf Kelu-1, in a search for rotational variability. Kelu-1 is a field brown dwarf, discovered by Ruiz, Leggett \\& Allard (1997) via its large proper motion. It is classified as an L2 dwarf in the scheme of Kirkpatrick et al. (2000), and model fits to its spectrum estimate an effective temperature of 1900$\\pm$100K (Ruiz et al.\\ 1997). The parallax distance of 19.6$\\pm$0.5\\,pc gives an absolute magnitude of M$_{\\textrm{\\small j}}$=11.96$\\pm$0.09 (Kirkpatrick et al.\\ 2000) or a bolometric magnitude of M$_{\\textrm{\\small bol}}$=13.9$\\pm$0.1. The detection of strong lithium absorption (Ruiz et al.\\ 1997) in an object of this luminosity implies a mass below 0.07M$_\\odot$ (Tinney 1998), making Kelu-1 a bona fide brown dwarf.  In common with many L- and M-type brown dwarfs (Basri et al.\\ 2000, Tinney \\& Reid 1998), Kelu-1 is a very rapid rotator with a measured $v \\sin i$=60$\\pm$5\\,kms$^{-1}$. It is thought this may indicate that the magnetic braking mechanisms which operate in more massive stars do not operate with the same efficiency in brown dwarfs (as first suggested by Tinney \\& Reid 1998).  In section 2 we describe the data acquisition and reduction. Time series analysis of the resulting differential photometry is presented in section 3. In section 4 we discuss our observations in terms of surface features on the brown dwarf, or the nature of an orbiting companion. ",
        "conclusions": "It is now clear that variability from brown dwarfs can be detected, but that photometry of better than 1\\% is required to do so. In the near future, futher studies of variability and rotation in brown dwarfs should greatly increase our understanding of their physics.  We have detected a strong periodicity of 1.80$\\pm$0.05 hours in differential photometry of the L2 brown dwarf Kelu-1. We have investigated four possible mechanisms to explain this variability:  \\begin{enumerate} \\item surface inhomogeneity moderated by meteorology and variable dust formation; \\item surface inhomogeneity moderated by magnetic starspots; \\item light curve variability due to gravitational distortion of Kelu-1's envelope by a close companion; and \\item light curve variability due to an eclipsing binary \\end{enumerate}  Mechanisms (ii) \\& (iv) seem unlikely explanations, but we are unable to concusively differentiate between mechanisms (i) \\& (iii). Ellipsodial variability, mechanism (iii), would produce a twin peaked lightcurve, giving a period of 3.6$\\pm$0.1 hours. This mechanism will give stable and repeatable photometric variability in future epochs. Alternatively, mechanism (i) associates the 1.80$\\pm$0.05 hour period with the rotation period, which is consistent with the rotational velocity of 60 kms$^{-1}$ and theoretical radius of $\\sim$0.1R$_\\odot$, indicating an inclination in the range 53\\deg $\\leq i \\leq$ 77\\deg.  Over the duration of our observations, the general shape and period of the lightcurve are unchanged (at least to within the measurement noise), implying the process causing the modulations is also stable on this timescale. This is in contrast to previous variability observations where no periodicity, or a period that changes on the order of the observation length, has been measured.  The two explanations we have presented lead to different predictions for future observations of Kelu-1. Variability induced by a secondary companion will be completely repeatable at future epochs, whereas long term evolution of surface features will result in a secular changes in the lightcurve. Further observations of Kelu-1 will be a powerful discriminant between the two hypotheses we have presented."
    },
    "0201/astro-ph0201481_arXiv.txt": {
        "abstract": "We present a method for solving the lightcurve of an eclipsing binary system which contains a Cepheid variable as one of its components as well as the solutions for three eclipsing Cepheids in the Large Magellanic Cloud (LMC).  A geometric model is constructed in which the component stars are assumed to be spherical and on circular orbits. The emergent system flux is computed as a function of time, with the intrinsic variations in temperature and radius of the Cepheid treated self-consistently.  Fitting the adopted model to photometric observations, incorporating data from multiple bandpasses, yields a single parameter set best describing the system.  This method is applied to three eclipsing Cepheid systems from the MACHO Project LMC database: MACHO ID's 6.6454.5, 78.6338.24 and 81.8997.87.  A best-fit value is obtained for each system's orbital period and inclination and for the relative radius, color and limb-darkening coefficients of each star.  Pulsation periods and parameterizations of the intrinsic color variations of the Cepheids are also obtained and the amplitude of the radial pulsation of each Cepheid is measured directly.  The system 6.6454.5 is found to contain a 4.97-day Cepheid, which cannot be definitely classified as Type I or Type II, with an unexpectedly brighter companion.  The system 78.6338.24 consists of a 17.7-day, W Vir Class Type II Cepheid with a smaller, dimmer companion.  The system 81.8997.87 contains an intermediate-mass, 2.03-day overtone Cepheid with a dimmer, red giant secondary. ",
        "introduction": "Large scale microlensing surveys have provided unprecedented resources for variable star research.  Their long time baseline and stable, accurate photometry are ideal for the detection and analysis of such objects and the large number of systems observed increases the probability of finding astrophysically-interesting objects that either have escaped detection or do not exist in our own galaxy.  Very few regularly-pulsating stars are known to belong to eclipsing systems.  One of the best studied is AB Cas, an Algol-type binary system, which contains a $\\delta$ Scuti type primary \\citep{abcas}. \\citet{R00} list 6 additional $\\delta$ Scutis which are members of eclipsing binaries.  The lone Galactic candidate for an eclipsing binary containing a Cepheid variable was BM Cas \\citep{Thiessen} but further study revealed that the variable was unlikely to be a Cepheid \\citep{fernev}.  The astrophysical benefits of a Cepheid variable in an eclipsing system could be considerable.  If the system is double-lined, a determination of the Cepheid's luminosity and mass can be made that is not only more accurate than existing measurements but also truly independent of the intervening steps in the distance ladder.  Such a system would provide the most direct measurement of the mass of a Cepheid and would offer an independent calibration of the period-luminosity relation.  Here we present the results of lightcurve analyses of three eclipsing Cepheid systems in the MACHO project Large Magellanic Cloud (LMC) database: MACHO ID's 6.6454.5, 78.6338.24 and 81.8997.87.  In Sec. 2 we describe the sources of the photometric observations.  Sec. 3 describes the model used to generate the lightcurve of an eclipsing Cepheid system and Sec. 4 describes the inverse problem of computing the parameters from an observed lightcurve.  In Sec. 5 we present the results obtained for the three systems.  Finally, Sec. 6 summarizes the analysis, describes work in progress and suggests future avenues for research. ",
        "conclusions": "The best fit parameters for each of the systems are shown in Tables \\ref{tab:6param}-\\ref{tab:81param}.  The orbital and pulsational periods are both in days and the inclination is given in degrees.  The radii (R) and amplitude of the radial variation ($\\Delta R_{amp}$) are relative to the orbital separation of the two stars.  $\\Delta R_{shift}$, which measures the shift of the radial change relative to the temperature change, has units of days.  Only the ratio of surface brightnesses $J_V/J_R$ is tabulated as it was the only one that was fit directly and the other surface brightnesses are computed from it as outlined above.  The limb-darkening coefficients, $x_\\lambda$ are as defined in equation \\ref{eq:ld}.  Once the relative surface brightness and radius of each star has been determined they can be combined with the mean system magnitude to compute the magnitude of each star in all filters.  Also computed for each of the three Cepheids is the value of $W_R = R - 4.0(V-R)$, an index which corrects for most of the effects of reddening and effective temperature differences.    Figures \\ref{fig:6vp}-\\ref{fig:81vp} show the primary eclipses for each system along with the best fit lightcurve.  Figure \\ref{fig:pl} shows a period-$W_R$ diagram (P-L diagram) for MACHO Cepheids and the locations of the three Cepheids studied here.  This diagram is essentially free of reddening and allows us to classify the three Cepheids under study based on their relation to other LMC Cepheids without an explicit correction for extinction.  The uncertainty ranges in the magnitude values in Tables \\ref{tab:6param}-\\ref{tab:81param} and Figure \\ref{fig:pl} were estimated from the range of possible component magnitudes based on the uncertainties in the best fit parameters.  These are statistical uncertainties and likely underestimate the true uncertainties.  The colors and magnitudes of each star allow some general comments on the evolutionary state of each system if we assume each star follows a standard, single-star evolutionary history.  This assumption is not unwarranted here given the large orbital periods of the systems. First, a crude correction for extinction must be made. To account for foreground reddening values of $E(B-V)$ are adopted from the map of Galactic foreground color excess toward the LMC published by \\citet{sch91}.  This yields values of $E(B-V)$ for the three systems as follows: 0.08 for 6.6454.5, 0.08 for 78.6338.24 and 0.10 for 81.8997.87.  These allow us to determine values for $A_V$ and $A_R$ when combined with the standard value of the ratio of total to selective extinction, $R_V = A_V/E(B-V) = 3.1$ \\citep{cousins}, and $A_R/A_V = 0.77$ for our Cousins $R$ and $V$ from the interstellar extinction relations of \\citet{cardelli}.  This procedure is clearly sufficient for the system 78.6338.24, producing $V_\\circ = 16.20 \\pm 0.03$ which agrees with the period-$M_v$ relation of \\citet{type2} for type II stars with log $P > 1.1$.  Applying the correction described above to the system 81.8997.87 fails to produce V and R magnitudes for the Cepheid that are consistent with those expected of an overtone Cepheid of its period.  Given the system's proximity to the 30 Doradus star-forming region it is not unreasonable to expect substantial extinction along this line of sight within the LMC and the closest Cepheid on the same plate, the overtone 81.8997.128, also appears well below the overtone band in the P-L diagram prior to applying a reddening correction.  The period-magnitude relation for overtone Cepheids of \\citet{baraffe} gives $V_\\circ = 15.86$ for the Cepheid in 81.8997.87, implying $A_v = 1.31$, a total value of $E(B-V) = 0.42$ and $R_\\circ = 15.56$.  Both $V_\\circ$ and $R_\\circ$ are consistent with the overtone bands in the period-V magnitude and period-R magnitude plots for MACHO LMC Cepheids.  A total value of $E(B-V) = 0.41$ was obtained by \\citet{demarchi} for selected regions of 30 Doradus.  A further correction may also be necessary for the system 6.6454.5, however, the ambiguity in the classification of the Cepheid (see below) precludes definitively comparing its properties to those expected from a period-luminosity relation as was done with the other two systems.  In light of this we adopt only the correction for foreground reddening.  After correction for reddening the values of $V$ and $V-R$ are converted to the $L - T_{eff}$ plane by assuming $\\mu_{LMC} = 18.5$ mag and using \\begin{equation} log(T_{eff}) = 4.199 - \\sqrt{0.08369+0.3493(V-R)} \\end{equation} a transformation of the \\citet{cwc} semi-empirical calibration. Figures \\ref{fig:iso1} and \\ref{fig:iso2} show theoretical isochrones representative of the metallicities of Type I and Type II stellar populations: Y=0.25, Z=0.008 isochrones from \\citet{bertelli} in Figure \\ref{fig:iso1} and Y=0.230, Z=0.0004 isochrones from \\citet{fagotto} in Figure \\ref{fig:iso2}, along with the properties of the three systems.  The error bars are computed exclusively from the errors in the magnitudes and colors given in Tables \\ref{tab:6param}-\\ref{tab:81param}.  Based on the properties tabulated in Table \\ref{tab:6param} the system 6.6454.5 is found to contain a Cepheid as the secondary with a brighter, bluer primary.  The nature of the Cepheid is unclear as its location in the P-L diagram (Figure \\ref{fig:pl}) places it between the Type I and Type II bands, inconsistent with both classifications. The large amplitude of the radial variation, 0.323 $\\pm$ 0.008 of the Cepheid's minimum radius is more consistent with a Type II classification.  Fundamental-mode Type I Cepheids have typical radial variations of of 10\\% or less \\citep{armstrong01} while Type II Cepheids show larger radial excursions in the range of 30-50\\% of the minimum radius \\citep{leb92}.  If the Cepheid is assumed to be a Type II Cepheid it is either making an excursion from the AGB or moving off the HB to the AGB (a less likely scenario given its period).  The companion, displayed in Figure \\ref{fig:iso2}, which also appears to be considerably evolved, is too luminous to fit either of these scenarios.  It is possible that this system is not in the LMC but is instead a foreground object.  A reduction in the assumed distance to the system to 17.8 kpc is necessary to shift the Cepheid's properties to fit a post-HB evolutionary state.  By contrast, if the system is compared to Type I isochrones (Figure \\ref{fig:iso1}) the Cepheid's location is consistent with that expected but the companion appears to be too blue to fit the isochrones.  If the Cepheid is indeed Type I, an additional reddening correction would need to be applied to make its location in Figure \\ref{fig:pl} consistent with the fundamental, Type I band. This would imply an even higher effective temperature and luminosity for the companion.  The status of 78.6338.24 is less ambiguous. It consists of a Type II Cepheid secondary with a hotter, but somewhat dimmer primary.  With a pulsational period in excess of 17.5 days the Cepheid would classified as a W Vir type, which is consistent with its large radius relative to its companion and with its location in the P-L diagram (Figure \\ref{fig:pl}).  This system presented several challenges to modelling. The pulsation period of the Cepheid was found to not be constant over the duration of the observations.  Despite being very small in magnitude (less than 1\\% of the period) this drift in period produced a substantial decrease in the quality of the fit.  It was corrected to some degree by assuming a pulsational angular frequency that was a slowly varying fuction of time, parameterized by: \\begin{equation} \\omega = \\omega_\\circ + A_1Bt + A_2(Bt)^2 \\label{eq:freq} \\end{equation} where the $A_i$ are parameters to be determined by fitting and B is a constant, set by trial and error, to ensure that the $A_i$ are of the same order as the other parameters.  Inspection of the complete set of residuals after fitting revealed indications of non-sphericity in one (or both) of the system components.  Both the asphericity and the ``period drift'' are consistent with the large radius and tenuous outer envelope of a W Vir star.  Their locations in Figure \\ref{fig:iso2} show both components to be well-evolved, post-HB or post-AGB objects.  The binary 81.8997.87 is distinct from the other two systems in several ways.  Its variable is an intermediate-mass Cepheid pulsating in the first-overtone mode.  Furthermore, the Cepheid is the primary with a considerably cooler, dimmer companion.  The amplitude of radial pulsations is low, 0.060 $\\pm$ 0.006 of the minimum Cepheid radius, as expected for an overtone Cepheid.  \\citet{gie82} provides an explicit determination of the radial displacement for the galactic overtone Cepheid SU Cas.  The radial amplitude of that star is 0.026 of the mean Cepheid radius, similar to the value we find for 81.8997.87 (note that workers tend to avoid Baade-Wesslink analyses of overtone Cepheids due to the small dynamic range of the observables).  Table 5a of \\citet{1978A&A....62...75P} lists physical properties derived from fitting model atmospheres to the continuum colors of a sample of Galactic Cepheids.  The values of $\\Delta R / \\langle R \\rangle$ obtained for the six overtones in the sample range from 0.042 to 0.080.  Once photometric contamination from the companion is removed the Cepheid's lightcurve shape, parameterized by the Fourier ratio $R_{21}$, is consistent with those of other MACHO overtone Cepheids of similar period.  Its location in the P-L diagram ({Figure \\ref{fig:pl}) further reinforces this view.  The companion's properties suggest a RGB star, possibly K class or later, however this combination of binary components is in poor agreement with current models of single star evolution.  This is reflected in the disagreement with the isochrones seen in Figure \\ref{fig:iso1}. Indeed any system consisting only of these two stars may be evolutionarily inconsistent.  The observational coverage of this system is far from ideal.  The 800-day orbital period limits the number of primary eclipses in the MACHO database to only 3.  Furthermore the secondary eclipses appear to be non-existent.  This can be accounted for by assuming a very low surface brightness for the secondary which will produce a very shallow secondary eclipse that could be dwarfed by the Cepheid variability. The absence of significant secondary eclipses could be a consequence of poor observational coverage.  The near 2-day pulsational period (2.03 days) combined with MACHO's single-point per night coverage results in repeated sampling of the same two pulsation phases during an individual eclipse.  These factors, poor coverage and the ill-defined secondary eclipses, produce large uncertainties in the fit parameters and component properties in Table \\ref{tab:81param} and Figure \\ref{fig:pl}.  In these eclipsing systems, the radial displacement of the Cepheid can be detected directly solely from modelling of the photometric lightcurve.  This is in contrast to most measurements of radial amplitude which are inferred from radial velocity measurements.  As a test, the fits of all three systems were repeated with the amplitude of the radial displacement held fixed at 0 and the resulting $\\chi^2$ values compared to those with radial variation included.  For the system 6.6454.5, $\\chi_v^2$ increased to 4.0, a 145\\% increase over the value of 1.6 listed in Table \\ref{tab:6param}.  For 81.8997.87 $\\chi_v^2$ increased to 1.5, an increase of only 26\\% over the best fit value of 1.1.  For 78.6338.24 $\\chi^2$ increased to 8.8 from 7.5, only a 17\\% change.  For each of the systems the change in $\\chi^2$ is found to be statistically significant for the number of degrees of freedom present.  The relatively smaller impact of the radial amplitude on the quality of the fit for 78.6338.24 could be explained by at least two factors: \\begin{enumerate}  \\item {In this system the eclipse duration and pulsation timescales are sufficiently similar that they produce a degeneracy in a parameter set which includes $\\Delta$R}.  \\item {The $\\chi^2$ value for this system is already elevated due to the model inadequacies previously mentioned.  Their impact on the fit could easily dwarf the effects of the inclusion of radial amplitude.}  \\end{enumerate}  For 81.8997.87 the small fractional amplitude of the radial change ($\\sim$6\\%) could make its effect on the lightcurve difficult to discern."
    },
    "0201/astro-ph0201448_arXiv.txt": {
        "abstract": "We present a theoretical analysis of galaxy-galaxy lensing in the context of halo models with CDM motivated dark matter profiles. The model enables us to separate between the central galactic and noncentral group/cluster contributions. We apply the model to the recent SDSS measurements with known redshifts and luminosities of the lenses. This allows one to accurately model the mass distribution of a local galaxy population around and above $L_{\\star}$.  We find that virial mass of $L_{\\star}$ galaxy is $M_{200}=(5-10)\\times 10^{11}h^{-1}M_{\\sun}$ depending on the color of the galaxy. This value varies significantly with galaxy morphology with $M_{\\star}$ for late types being a factor of 10 lower in $u'$, 7 in $g'$ and a factor of 2.5-3 lower in $r'$, $i'$ and $z'$ relative to early types. Fraction of noncentral galaxies in groups and clusters is estimated to be below 10\\% for late types and around 30\\% for early types. Using the luminosity dependence of the signal we find that for early types the virial halo mass $M$ scales with luminosity as $M \\propto L^{1.4 \\pm 0.2}$ in red bands above  $L_{\\star}$. This shows that the virial mass to light ratio is increasing with luminosity for galaxies above $L_{\\star}$, as predicted by theoretical models. The virial mass to light ratio in $i'$ band is $17(45)hM_{\\sun}/L_{\\sun}$ at $L_{\\star}$ for late (early) types. Combining this result with cosmological baryon fraction one finds that 70(25)\\%$h^{-1}\\Upsilon_i\\Omega_m/12\\Omega_b$ of baryons within $r_{200}$ are converted to stars at $L_{\\star}$, where $\\Upsilon_i$ is the stellar mass to light ratio in $i'$ band. This indicates that both for early and late type galaxies around $L_{\\star}$ a significant fraction of all the baryons in the halo is transformed into stars. ",
        "introduction": "Weak lensing by matter along the line of sight between the source and the observer shears the images of the background galaxies, inducing ellipticity distortions (see \\citeNP{2001PhR...340..291B} for a review of weak lensing). Although away from rich clusters the effect is too small to be detectable for individual galaxy lenses, it can be measured statistically as a function of relative separation from the galaxy. This requires averaging over the tangential ellipticities of all the background galaxies relative to the lens and over all the lenses (galaxies which are in the foreground). Until recently this averaging, named galaxy-galaxy (g-g) lensing, was done as a function of apparent angular position in the sky (\\shortciteNP{1984ApJ...281L..59T}, \\citeNP{1996ApJ...466..623B}, \\shortciteNP{1998ApJ...503..531H}, \\shortciteNP{2000AJ....120.1198F}, \\citeNP{2001ApJ...551..643S}), so a signal at a given angular separation could be either coming from a small radial distance of a nearby lens or from a large distance of a far lens. This made the theoretical interpretation of the data rather involved. First attempt to use distances was by \\shortciteN{2001ApJ...555..572W}, which however only had limited photometric information and so could only obtain reliable distances to early type galaxies.  Recent study of galaxy-galaxy lensing by SDSS collaboration \\shortcite{2001astro.ph..8013M} is a significant step forward in the study of galaxy-galaxy lensing. The spectroscopic sample of more than 35,000 lensing galaxies and 3.6 million background galaxies is large enough to allow one a detailed study of the relation between mass and light for several luminosity bands and morphological types. Since the distances for lens galaxies are known one can study the strength of the signal as a function of proper radial separation from the lens. In addition, because the survey is shallow, the redshift distribution of background galaxies is known from the deeper spectroscopic surveys (e.g. \\citeNP{1995ApJ...455...50L}). In combination with above this means that the mean critical density is known for every lens, so one can average over the proper projected mass density rather than the shear itself. This fact greatly simplifies the theoretical analysis, since one can now measure the actual projected density as a function of galaxy luminosity and proper radial transverse distance from the galaxy. In addition, since the lens sample is at low redshift (mean $\\bar{z}\\sim 0.1$) redshift evolution is small and k-corrections are relatively reliable in red bands, further simplifying the theoretical interpretation.  In this paper we want to connect the SDSS observational results to the theoretical models in the context of our current understanding of galaxy formation models within the CDM paradigm. We will model the dark matter halos with CDM type of halo profiles (NFW profile; \\citeNP{1997ApJ...490..493N}), where the slope is gradually changing from the inner slope between -1 to -1.5 to the outer slope of -3. Although other profiles have been proposed that differ significantly from NFW in the inner parts of the halo, they agree well with NFW in the outer parts \\shortcite{2001ApJ...554..903K}. NFW profile should be contrasted to the truncated singular isothermal sphere with a constant slope -2 out to a fixed radius which was often used used in the past work on g-g lensing. Our main goal is to determine the virial mass of the halo and its relation to the luminosity of the galaxy. This is important for theoretical models of galaxy formation, since it is usually assumed that only baryons within the virial radius are able to condense and form stars. By determining the virial mass for a given galaxy luminosity one can thus directly determine the efficiency of star formation in typical galactic halos. Luminosity dependence of the galaxy-galaxy lensing signal allows one to determine this as a function of halo mass, while morphology subsamples can determine it as a function of morphological type. By comparing the radial dependence of the signal with the theoretical models one can also determine the fraction of galaxies that live in larger halos such as groups and clusters. This is another parameter that can distinguish between the different galaxy formation models. Finally, comparison of high density sample to the field sample allows one to study the effects of dense enviroments, such as tidal stripping, on the dark matter profile and mass to light ratios. Thus g-g lensing allows one to make detailed tests of the galaxy formation models (e.g. \\shortciteNP{1999MNRAS.303..188K}, \\shortciteNP{2000MNRAS.311..793B}, \\citeNP{1999MNRAS.310.1087S}). In principle g-g lensing can also allow a direct determination of the dark matter halo profile, although as we will show here at present the data do not have enough power to strongly constrain this.  An important issue in the theoretical analysis of g-g lensing is how to separate the contribution from the individual halo of the galaxy from that of the neighbouring galaxies or larger mass concentrations such as groups and clusters. The former should dominate on small scales while the latter on large scales. Previous work used the additional information obtained from the galaxy clustering to remove this contribution, but in general such analysis relies on the assumption that all the mass is associated with galaxies, which is invalid on scales where groups and clusters become dominant. Most of the mass in these systems is in a diffuse form and only about 5-10\\% is expected to be attached to individual galaxies (e.g. \\shortciteNP{2001MNRAS.328..726S}, \\shortciteNP{2000ApJ...544..616G}, \\shortciteNP{2001MNRAS.321..559B}). This means that these systems cannot be modeled using the galaxies as the mass tracers, which only accounts for a small fraction of the total mass. In fact, theoretical models presented in this paper suggest the contribution from these group and cluster halos dominates the signal on scales above 200$h^{-1}$kpc (see also \\citeNP{2000MNRAS.318..203S}) and has to be carefully modeled to account for it.  Given that corrections to the profile on large scales are difficult to determine from the data directly one has to turn to theory for guidance. Realistic theoretical modeling of g-g lensing must combine galaxy formation models and dark matter models. One way is to use semi-analytic or hydrodynamic models of galaxy formation and combine them with N-body simulations (\\citeNP{2001MNRAS.321..439G}; \\citeNP{2001astro.ph..7023W}). This has the advantage of having a realistic distribution of dark matter and galaxies, but suffers from the lack of force and mass resolution. In addition, this approach by itself does not allow for a fast exploration of parameter space. The alternative approach is to use the recently popularized halo model applied to galaxies (\\citeNP{2000MNRAS.318..203S}, \\citeNP{2000MNRAS.318.1144P}, \\shortciteNP{2001ApJ...546...20S}, \\citeNP{2001astro.ph..9001B}), which takes into account both the individual halo profiles (for galaxies either at the centers of the galactic halos or distributed within larger groups and clusters with a specified radial distribution) and correlations between the galaxies. This approach applied to the galaxy-dark matter correlations (which fully determines the g-g lensing signal) was shown to give the same results as the simulations in the regime of applicability \\cite{2000MNRAS.318..203S}. However, halo model is analytical, does not suffer from the resolution issues and provides a more physical interpretation of the results. Thus one can parametrize the model with the quantities one wishes to extract from the data and determine these directly. It also allows for a rapid exploration of the parameter space without the need to rerun cosmological simulations or to repopulate the halos with galaxies using semi-analytic galaxy formation models. In this sense the halo model provides a natural link between the observations and the theoretical models of galaxy formation.  Even though we will use the halo model to analyze g-g lensing the main features can be understood without it. Particularly robust conclusions are possible using the low density sample, where the clustering and group/cluster contribution can be neglected. In this case the signal can be taken simply as a projected radial mass profile of the halo averaged over the mass distribution of the halos determined by the galaxy sample. For the full and high density sample the contribution from groups and clusters can no longer be neglected, as it dominates the signal on scales above 200$h^{-1}$kpc. This allows one to determine the fraction of galaxies in groups and clusters as well.  The outline of this paper is as follows: in section \\S2 we review the basic theory and halo model as applied to g-g lensing. In \\S3 we discuss the influence of the parameters introduced on the observed g-g lensing signal. We focus on the relative contributions to g-g lensing from galaxy, group/cluster and clustering terms and explore which parameters the observations are most sensitive to. In \\S4 we apply the model to the data to determine several of the model parameters. Interpretation of the results and conclusions are presented in \\S5. ",
        "conclusions": "This paper establishes a quantitative framework on how to analyze g-g lensing data and applies it to the SDSS sample of 35,000 galaxies with known redshifts and luminosities. The main qualitative new feature of the model is that both the dark matter around the individual galaxies and dark matter in groups and clusters are included. Both components are needed to explain the observations. The two contributions have a different radial dependence and can be determined separately. This provides important constraints on the galaxy formation models, which must satisfy the relative contribution from central and group/cluster components. We also argue that correlations between the galaxy and another halo that is not the one that galaxy belongs to can be neglected, at least on scales below 1$h^{-1}$Mpc of interest here.  The main result of this paper is determination of galaxy virial masses and the fraction of galaxies in groups and clusters as a function of luminosity and morphology. These provide important constraints on the galaxy formation models. The average virial mass $M_{200}$ of an $L_{\\star}$ galaxy is around $(5-10)\\times 10^{11}h^{-1}M_{\\sun}$, depending on the passband. This mass varies significantly with morphology and the variation is largest in $u'$, where the difference between early and late types can be up to a factor of 10, decreasing to a factor of 2-3 in $r'$, $i'$ and $z'$. For example, in $i'$ we find $M_{200}=(3.4 \\pm 2.1) \\times 10^{11}h^{-1}M_{\\sun}$ for late types and $M_{200}=(9.3\\pm 2.2) \\times 10^{11}h^{-1}M_{\\sun}$ for early types. While the signal for late types is rather weak and consequently the errors are large, they are gaussian distributed and so we can exclude the possibility that in red bands the $M_{\\star}$ for early and late types are equal.  What does this imply for the star formation efficiency as a function of galaxy morphology? To address this we first transform from $i'$ to K band, where the luminosity is assumed to be a reliable tracer of stellar mass for a given age of the population (and IMF). This transformation does not change the results significantly, since $K-i'$ differs by 0.2-0.3 magnitudes at most between the two types \\shortcite{2001astro.ph.11024I}. This difference is further reduced by up to 20\\% because of the missed light for the de-Vaucouleurs profile. As a result, the difference in $M/L$ between the two types is reduced to a factor of 2. In K band the difference in $M_{\\rm stellar}/L$ between early and late type population can be up to a factor of 2 depending on the exact ages and IMF assumed, so our results are consistent with the assumption that star formation efficiency for early and late type galaxies is the same. The errors however are still large both on the virial masses and stellar mass to light ratios and there could still be up to a factor of 2 difference. Since the current results are based on only 5\\% of final SDSS sample one should be able to place much better constraints on this issue in the future.  If we adopt the late(early) type stellar mass to light ratio $M_{\\rm stellar}/L \\sim 1.5(3)h$ in $i'$ we find that at $L_{\\star}$ about 10-15\\% of virial mass is converted to stars. The fraction of baryons inside a halo should equal $\\Omega_b/\\Omega_m$, which can vary between 10-20\\% for $\\Omega_m \\sim 0.3 \\pm 0.1$ and $\\Omega_b \\sim 0.04 \\pm 0.01$. We see that if our halo masses are correct a significant fraction of baryons, up to 100\\%, is converted to stars in such halos. For early types this fraction decreases with luminosity and is a factor of 2 lower at $7L_{\\star}$. These stellar fractions can be accomodated in the standard models, but require a low matter density and/or a high baryon density, so that the overall baryon to dark matter fraction is sufficiently high.  We find that $M/L \\propto L^{0.4\\pm 0.2}$ in red bands ($i'$ and $z'$). An increase in $M/L$ with luminosity above $L_{\\star}$ is expected theoretically from semi-analytic models of galaxy formation (\\shortciteNP{1999MNRAS.303..188K}, \\shortciteNP{2000MNRAS.311..793B}). Comparison of these results with mass tracers such as Tully-Fisher or fundamental plane is complicated, because these probe the halo at smaller radii, where the rotation velocity of dark matter could be different and where baryons could play a significant role. A detailed analysis will be presented elsewhere. The results from such analysis show that the masses obtained here are consistent with the CDM picture in which the rotation velocity drops by roughly 1.7-1.8 from the optical to the virial radius at $L_{\\star}$. Such a drop is expected in CDM models, where the rotation curve peaks at a fraction of the virial radius and drops beyond that. However, to explain such a large drop one also needs a significant additional contribution to the rotation velocity from the stellar component and/or baryon compressed dark matter. For early types the scaling $M\\propto L^{1.5}$ indicates $L \\propto v_{200}^{2}$, while for the same galaxies the Faber-Jackson relation gives $L \\propto \\sigma^{4}$ \\cite{2001astro.ph.10344B}. In this case the ratio of optical to virial circular velocity depends on luminosity and drops to 1.4 at $7L_{\\star}$.  Comparison with other g-g lensing results is also complicated, since none of these use luminosity information, are typically at a higher redshift with poorly determined distances and may have a different morphological composition. The most direct comparison can be made to the study of early type galaxies by \\shortciteN{2001ApJ...555..572W}, where color information was used to determine approximate redshifts to these galaxies. A direct comparison shows that their virial mass for early type $L_{\\star}$ galaxy is around $2\\times 10^{12}h^{-1}M_{\\sun}$. Their sample is at a higher redshift and may not be directly comparable to ours, but the obtained value is quite close to our best fitted value of $M_{\\star}=2.1\\times 10^{12}h^{-1}M_{\\sun}$ in $g'$, which is the closest to $B$ band used there (note that $L_{\\star}=1.1\\times 10^{10}h^{-2} L_{\\sun}$ in both samples). In their analysis they assume $\\beta=0.5$, which is inconsistent with the SDSS luminosity dependent data, so their actual value of $M_{\\star}$ could even be somewhat lower because $\\beta$ and $M_{\\star}$ are anticorrelated. Earlier analysis of the dynamics of satellites around spiral galaxies gave significantly higher masses (up to a factor of 5, \\shortciteNP{1997ApJ...478...39Z}) and also did not show correlation between light and mass, but a more recent analysis of the SDSS data shows a better agreement with g-g lensing results (T. McKay, private communication). At the upper end of the luminosity range our results can be compared to group and cluster velocity dispersion analysis of \\shortciteN{2001astro.ph.12534G}. Our results agree well both on the mean luminosity of a $10^{13}h^{-1}M_{\\sun}$ halo and on the scaling of mass with luminosity. For late type galaxies at $L_{\\star}$ our results agree well with the virial mass derived from theoretical models by \\citeN{2001astro.ph.12566V}.  There are many aspects of the analysis that could be improved upon with better statistics which will be available in the future. We mention some here. We have assumed a power law relation between mass and luminosity with a constant slope $\\beta$, while theoretical models suggest $\\beta$ changes with mass. With the current sample the signal for galaxies less luminous than $L_{\\star}$ is too small to be detectable and the relation between mass and luminosity cannot be established in that range. As larger samples become available this regime will be probed as well. Fraction of galaxies in groups and clusters should be determined as a function of luminosity. Similarly, galaxy morphology dependent analysis should be improved by analyzing luminosity dependence of the signal for each morphology type and by dividing the lens galaxies into several morphological classes. With better statistics the fraction of galaxies in groups and clusters could be determined as a function of luminosity and not just for the overall sample as done here. The radial distributions of dark matter and galaxies in groups and clusters as well as galaxy occupation statistics as a function of halo mass should be quantified better. These effects can change the values for $M_{\\star}$ and $\\beta$ by 20-30\\% and remain the dominant source of systematic error. Better k-correction and evolution effects should be applied to the data to match the bright end (which receives contributions from  $z \\sim 0.2$) and faint end (low $z$) samples. This effect should not exceed 0.1-0.2 magnitudes in red bands and so should not change our conclusion on the mass luminosity relation significantly. Differences between halos of galaxies in the field and in groups and clusters could be studied in more detail, testing for example the collisionless CDM paradigm. In this paper we assumed that the profiles of galaxies within groups and halos are unchanged out to the truncation radius. In the extreme opposite case where galaxies inside groups and clusters do not retain any mass at all the virial masses of galaxies in the field will be underestimated by $1-\\alpha$. This is a small correction for late type galaxies, but potentially more important for early type galaxies. Finally, the clustering correction, which corrects for the fraction of galaxies that are correlated with the lens and are not in the background, is luminosity dependent and should be done more accurately. While the procedure adopted here is reasonable, it should be studied in more detail. This effect may be removed for example if one can select background galaxies using their photometric redshift (photo z) information, although it is not clear if the faint galaxies that dominate the background population have sufficient signal to noise for this purpose. Current experience with photo z's indicates that in SDSS these work only down to $r'<20.5$, while most of the background galaxies used for shape information are fainter than that.  As we obtain more data better statistical analysis will also be needed. Rather than divide the data in arbitrary luminosity bins, which may still be too broad for a quantitative analysis, one should parametrize the model with a few parameters and use maximum likelihood type of analysis to determine these. Photometric redshifts would also help to improve the statistics since one could weight the signal by giving optimal weight for each lens-background galaxy pair (this would then automatically downweight all the galaxies that are close to the lens). These improvements will allow one to make very robust statements on the relation between the virial mass and luminosity and on the membership in groups and clusters over a broad range of luminosities and morphological types.  US acknowledges he support of NASA, David and Lucille Packard Foundation and Alfred P. Sloan Foundation. J.G. was supported by grants 5P03D01820 and 2P03D01417 from Polish State Committee for Scientific Research. We thank Tim McKay, Erin Sheldon and Iskra Strateva for help with the interpretation of SDSS data and its analysis."
    },
    "0201/astro-ph0201413_arXiv.txt": {
        "abstract": "{A detailed study of the planetary nebula NGC~6565 has been carried out on long-slit echellograms ($\\lambda$/$\\Delta\\lambda$=60000, spectral range=$\\lambda\\lambda$3900--7750\\AA) at six, equally spaced  position angles. The expansion velocity field, the c(H$\\beta$) distribution and the radial profile of the physical conditions (electron temperature and density) are obtained. The distance, radius, mass and filling factor of the nebula and the temperature and luminosity of the central star are derived. The radial ionization structure is analyzed using both the classical method and the photo-ionization code CLOUDY.  Moreover, we present the spatial structure in a series of images from different directions, allowing the reader to ``see'' the nebula in 3-D. NGC~6565 results to be a young (2000--2500 years), patchy, optically thick triaxial ellipsoid (a=10.1 arcsec, a/b=1.4, a/c=1.7) projected almost pole-on. The matter close to major axis was swept-up by some accelerating agent (fast wind? ionization? magnetic fields?), forming two faint and asymmetric polar cups.  A large cocoon of almost neutral gas completely embeds the ionized nebula. NGC~6565 is in a recombination phase, because of the luminosity drop of the massive powering star, which is reaching the white dwarf domain (logT$_*\\simeq$5.08 K; logL$_*$/L$_\\odot\\simeq$2.0). The stellar decline started about 1000 years ago, but the main nebula remained optically thin for other 600 years before the recombination phase occurred.  In the near future the ionization front will re-grow, since the dilution factor due to the expansion will prevail on the slower and slower stellar decline. NGC~6565 is at a distance of 2.0($\\pm$0.5) Kpc and can be divided into three radial zones: the ``fully ionized'' one, extending up to 0.029--0.035 pc at the equator (0.050 pc at the poles), the ``transition'' one, up to 0.048--0.054 pc (0.080 pc), the ``halo'', detectable up to 0.110 pc.  The ionized mass ($\\simeq$0.03 M$_\\odot$) is only a fraction of the total mass ($\\ge$ 0.15 M$_\\odot$), which has been ejected by an equatorial enhanced superwind of 4($\\pm$2)$\\times$10$^{-5}$ M$_\\odot$ yr$^{-1}$ lasted for 4($\\pm$2)$\\times$10$^3$ years. ",
        "introduction": "\\begin{figure*} \\centering \\caption{WFPC2 appearance of NGC~6565 in [OIII] (left) and [NII] (right). The images have been retrieved from the HST archive. North is up and East is to the left.} \\end{figure*}  The late evolution of low and intermediate mass stars (1.0 M$_\\odot<$M$_*<$8.0 M$_\\odot$) is characterized by the planetary nebula (PN) metamorphosis: an asymptotic giant branch (AGB) star gently pushes out the surface layers in a florilegium of forms, crosses the HR diagram and reaches the white dwarf regime (Aller 1984, Pottasch 1984, Osterbrock 1989).  So far the interpretation of the exuberant morphologies exhibited by the PNe in terms of detailed three-dimensional structures and physical conditions of the ionized gas was limited by projection effects, leading to approximate spatial forms and to unrealistic assumptions for the main parameters, e.g. electron temperature and electron density constant all over the nebula (Aller 1984, 1990, 1994).  A 3-D reconstruction technique for studying at large and small scales the morphology, physical conditions, ionization, spatial structure and evolutionary status of PNe has been introduced by Sabbadin et al. (2000a, b) and Ragazzoni et al. (2001), based on echellograms of moderate spectral resolution (R$\\sim22000-25000$). The key of the 3-D methodology is simple: the PN is an expanding plasma. Thus the position, thickness and density of each elementary volume can be derived from the radial velocity, width and flux of the corresponding emission.  The procedure has been applied to:  \\begin{description} \\item[-] NGC~40: an optically thick, very low excitation barrel-shaped nebula with thin arcs emerging at both ends of the major axis, powered by a luminous and ``cold'' WC8 star presenting a large mass-loss rate. The fast, hydrogen depleted photospheric material ejected by the nucleus is gradually modifying the chemical composition of the innermost nebular regions (Sabbadin et al. 2000a); \\item[-] NGC~1501: an evolved, high excitation, optically thin oblate ellipsoid, denser in the equatorial belt, deformed by several bumps, embedded in a homogeneous, inwards extended cocoon and ionized by a `` hot'' and luminous WC4 star exhibiting nonradial g-mode pulsations (Sabbadin et al. 2000b; Ragazzoni et al. 2001). \\end{description}  In order to deepen the analysis, we have started a survey at high spectral and spatial resolutions with the ESO NTT. The superb quality of this material allows us to study at unprecedented accuracies objects with different morphology, e.g. NGC~7009 (the \"Saturn\" nebula), the tetra-lobed IC~4634, the butterfly HB~5, Mz~3 and NGC~6537, the double-envelope NGC~5882, NGC~6153 and NGC~6818.  To simplify the application of the 3-D method we decided to begin with an \"easy\" nebula without FLIERS (fast, low ionization emitting regions), BRETS (bipolar, rotating, episodic jets), ansae, wings, multiple envelopes etc., and naively selected NGC~6565.  \\begin{figure} \\caption{[NII]/[OIII] distribution over NGC~6565 (original frames and orientation as in Fig. 1), showing the large stratification effects and the faint, low ionization regions protruding from the main nebula in PA$\\simeq$145$\\degr$.} \\end{figure} ",
        "conclusions": "The long travel dedicated to the study of NGC~6565 is over. We started with the kinematical properties, passed through the physical conditions and arrived to the distance and radius of the nebula and to the temperature and luminosity of the star. The radial ionization structure was later explored, using both the ``classical'' method and the photo-ionization code CLOUDY, and the spatial structure investigated by assembling the tomographic maps.  In nuce: NGC~6565 is a young (2000--2500 yr), patchy, optically thick ellipsoid with extended polar cups, seen almost pole-on. It is in a recombination phase, because of the luminosity drop of the massive powering star, which is reaching the white dwarf domain. The stellar decline started about 1000 years ago, but the nebula remained optically thin for other 600 years before the recombination phase occurred (thus, at Galileo Galilei times NGC~6565 was at high excitation, a bit larger and much brighter than at present). In the near future the ionization front will re-grow, since the dilution factor due to the expansion will prevail on the slower and slower luminosity decline.  The nebula, at a distance of 2.0 Kpc, can be divided into three radial zones:  - the ``fully ionized'', extending up to 0.029--0.035 pc at the equator ($\\simeq$0.050 pc at the poles),  - the ``transition'', up to 0.048--0.054 pc ($\\simeq$0.080 pc at the poles),  - the recombining ``halo'', detectable up to 0.110 pc.  The ionized mass ($\\simeq$0.03 M$_\\odot$) is only a fraction of the total mass (M$_{\\rm tot}$$\\ge$0.15 M$_\\odot$), which has been ejected by an equatorial enhanced superwind of 4($\\pm$2)$\\times$10$^{-5}$ M$_\\odot$ yr$^{-1}$ lasted for 4($\\pm$2)$\\times$10$^3$ years.  From all the points of view NGC~6565 results to be an exciting laboratory which deserves further attention. Amongst the many and important aspects left unresolved by this paper we mention:  \\begin{description} \\item[-] the neutral gas distribution and kinematics. Detailed radio and infrared observations are highly advisable; for instance, the spectroscopically resolved structure in the H$_2$ emission at 2.122 $\\mu$m, which is a signpost of the bipolar structure (Kastner et al. 1996); \\item[-] the filling factor. We obtain $\\epsilon_{\\rm l}$=0.25 in the external, low ionization regions by comparing Ne(H$\\alpha$) and Ne[SII]. The same analysis should be performed in the internal, high ionization parts using  [ArIV] ($\\lambda$4711\\AA/$\\lambda$4740\\AA) and [ClIII] ($\\lambda$5517\\AA/$\\lambda$5537\\AA). Deeper echellograms are required; \\item[-]  the faint and fast (shocked?) polar cups, overlooked in our study mainly focused on the equatorial regions; \\item[-] the mechanisms and physical processes forming and shaping a PN like NGC 6565. This point appears problematic, mainly because of the huge number of models proposed for bipolar PNe.  They include: single AGB progenitors (+ fast wind + magnetic fields + photoionization), a planetary system, a binary system undergoing common envelope evolution, a close companion presenting a dense accretion disk (for details, see Frank 1999 and Soker \\& Rappaport 2001).  In all these models the nebular shaping is a slow and gradual process occurring in a time scale comparable to the PN life, whereas the superior spatial resolution of the HST imagery has recently revealed that bipolarity is a common feature already in proto-PNe (see Sahai \\& Trauger 1998, and references therein). This, on the one hand means that some imprint agent acts in the late AGB and/or early post-AGB phase, driving most of the subsequent nebular evolution, on the other hand stresses the inadequacy of the current models. \\end{description}  In addition to the specific analysis of our nebula, the aim of the present work was the search for a satisfying reduction procedure allowing us to exploit the huge amount of information contained in the high resolution spectra. They essentially cover two wide fields: kinematics (including expansion velocity, tomography and spatial structure), and physical conditions (radial profile of Ne, Te, ionization, chemical abundances etc.). The difficulties connected to the determination of the different parameters are perfectly synthesized in the Aller's (1994) sentence: ''A nebula is a three-dimensional structure for which we obtain a two-dimensional projection''. With the introduction of the 3-D methodology, the second half of the sentence becomes: ``... for which we obtain the three-dimensional structure''.  The accuracy of the 3-D reconstruction is defined by the ``relative'' spectral and spatial resolutions (Ragazzoni et al. 2001). The first is given by RR=$\\Delta$V/Vexp, $\\Delta$V being the spectral resolution. RR mainly depends on the intrumentation, since the expansion velocity range of the PNe is quite sharp (indicatively: 20 km s$^{-1}$$<$Vexp$<$40 km s$^{-1}$). The relative spatial resolution, SS=d/$\\Delta$d (d=angular extent, $\\Delta$d=seeing+guiding), essentially depends on the target, due to the large spread of apparent sizes exhibited by the PNe. In the case of a compact object like NGC~6565, the ESO NTT+EMMI echellograms give RR$\\simeq$5 and SS$\\simeq$10; a significant increase of the spatial resolution is expected for the largest nebulae of the sample (like NGC~7009 and NGC~6818).  In conclusion, we believe that the crucial point for the physical interpretation of the PNe is the detailed knowledge of the radial density profile, now obtainable from the H$\\alpha$ flux distribution in the zero-velocity pixel column. It opens the possibility of creating a realistic model for each expanding nebula (this includes: PNe, nova and supernova remnants, shells around Population I Wolf-Rayet stars, nebulae ejected by symbiotic stars, bubbles surrounding early spectral type main sequence stars etc.)."
    },
    "0201/astro-ph0201280_arXiv.txt": {
        "abstract": "FUSE high resolution spectra of two PG1159 type central stars (K1-16 and NGC\\,7094) have revealed an unexpected iron deficiency of at least 1 or 2 dex (Miksa \\etal 2002). Here we present early results of FUSE spectroscopy of the CSPN Abell~78. It is shown that iron is strongly deficient in this star, too. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0201/astro-ph0201249_arXiv.txt": {
        "abstract": "{ We describe an automated search through the Leiden/Dwingeloo \\hi Survey (LDS) for high--velocity clouds north of $\\delta=-28^\\circ$. From the general catalog we extract a sample of isolated high--velocity clouds, CHVCs: anomalous--velocity \\hi clouds which are {\\it sharply bounded in angular extent} with no kinematic or spatial connection to other \\hi features down to a limiting column density of 1.5$\\times$10$^{18}$ cm$^{-2}$. This column density is an order of magnitude lower than the critical \\hi column density, $\\sim$2$\\times$10$^{19}$ cm$^{-2}$, where the ionized fraction is thought to increase dramatically due to the extragalactic radiation field. As such, these objects are likely to provide their own shielding to ionizing radiation. Their small median angular size, of about $1^\\circ$ FWHM, might then imply substantial distances, since the partially ionized \\hi skin in a power--law ionizing photon field has a typical exponential scale--length of 1~kpc. The automated search algorithm has been applied to the HIPASS and to the Leiden/Dwingeloo data sets.  The results from the LDS are described here; Putman et al. (\\cite{putman02}) describe application of this algorithm to the HIPASS material.  We identify 67 CHVCs in the LDS which satisfy stringent requirements on isolation, and an additional 49 objects which satisfy somewhat less stringent requirements.  Independent confirmation is available for all of these objects, either from earlier data in the literature or from new observations made with the Westerbork Synthesis Radio Telescope and reported here.  The catalog includes 54 of the 65 CHVCs listed by Braun \\& Burton (\\cite{braun99}) on the basis of a visual search of the LDS data.  \\keywords ISM: clouds -- ISM: kinematics and dynamics -- Galaxy: evolution -- Galaxies: dwarf -- Galaxies: evolution -- Galaxies: Local Group } ",
        "introduction": "High--velocity clouds (HVCs) were first encountered in the $\\lambda\\,21$ cm line of \\hi at radial velocities unexplained by any conventional model of Galactic rotation.  Since their discovery by Muller et al.~(\\cite{muller63}), they have remained enigmatic objects of continued interest.  Wakker \\& Van Woerden~(\\cite{wakker97}) and Wakker et al.~(\\cite{wakker99}) have given recent reviews; since these reviews, progress has been made on several fronts.  The anomalous--velocity clouds are found scattered over the entire sky, and examples are found throughout a range of radial velocity spanning about 800 \\kms: obtaining an adequate observational foundation for the phenomenon has been a persistent and continuing challenge.  We describe here a search algorithm which has been applied to the all--sky coverage afforded by the new \\hi surveys of the northern and southern skys, and the results of its application to the Leiden/Dwingeloo Survey for examples of the phenomenon.  During the past forty years, a wide variety of explanations for the HVCs has been suggested.  The matter of distances has remained particularly difficult, and it is on distances that most of the basic physical properties depend. Only in a few cases have distances been measured or constrained.  The distinct system of anomalous--velocity features recognised as the Magellanic Stream represents tidal debris originating in a gravitational interaction of the Large and Small Magellanic Clouds with our Galaxy (see Putman \\& Gibson \\cite{putman99}), and is therefore likely to be located at distances of several tens of kpc.  Other distinct systems of high--velocity objects constitute a few complexes, stretching over regions of some tens of square degrees.  One of these, Complex A, has been found from absorption--line observations (van Woerden et al. \\cite{vanwoerden99}; Wakker \\cite{wakker01}) to lie within the distance range $8 < d < 10$ kpc.  But the term HVC has been used to encompass a wide range of phenomena; unlike the Magellanic Stream and the half--dozen well--known complexes, many of the individual anomalous--velocity features are compact and are isolated on the sky down to low column density limits, as we will demonstrate below.  The properties of anomalous--velocity \\hi emission might be more readily determined after a classification into sub--categories has been made.  After compiling a general catalog of high--velocity features in the northern sky, we focus on identifying the category of compact, isolated features, which show no connection in position and velocity with the Galaxy, the Magellanic Clouds or the extended HVC complexes.  Braun \\& Burton (\\cite{braun99}) have argued that these objects may represent a single class of clouds, whose members originated under similar circumstances and which share a common evolutionary history, and which might lie scattered throughout the Local Group.  The idea that the anomalous--velocity clouds are deployed throughout the Local Group has been considered earlier, by (among others) Oort (\\cite{oort66}, \\cite{oort70}, and~\\cite{oort81}), Verschuur~(\\cite{verschuur75}), Eichler~(\\cite{eichler76}), Einasto et al. (\\cite{einasto76}), Giovanelli (\\cite{giovanelli81}), Arp~(\\cite{arp85}), and Bajaja et al. (\\cite{bajaja87}).  Various arguments have been raised against these interpretations.  In the first review of the possible interpretations of high--velocity clouds, Oort (\\cite{oort66}) ruled out the supposition that the clouds could be independent systems in the Local Group on two principal grounds: he stated that `` ... a situation outside our Galaxy would give no explanation of the principal characteristic of the high--latitude clouds, viz. that the high velocities ... are all negative\", and furthermore that `` ... it would be almost impossible to explain on this hypothesis high--velocity clouds which apear to be related with each other over regions $30\\deg$ or more in diameter\".  Since Oort's first review, newer \\hi surveys have extended the sky coverage and have revealed that there are, in fact, approximately as many (compact) anomalous--velocity clouds at positive velocities as there are at negative velocities; and the objection against the large angular size of the complexes is confronted by the knowledge that these features, in any case, are indeed located within the Galactic halo.  By analyzing the stability of a median HVC (in the Wakker \\& Van Woerden \\cite{wakker91} tabulation) against Galactic tidal disruption and self--gravity, Blitz et al.~(\\cite{blitz99}) suggest a distance of~1~Mpc, for an assumed ratio between \\hi mass and total mass of~0.1. Furthermore they suggest that the preferred coordinate system for the clouds is neither the Local Standard of Rest system, nor the Galactic Standard of Rest system, but the Local Group Standard of Rest system. The amplitude of the average velocity and the velocity dispersion of the cloud system both have the lowest values in this system, indicating that it might be the most relevant.  A numerical simulation of the dynamics of a population of low mass test masses within the gravitational potential of the Milky Way and M~31, reproduces some aspects of the kinematic and spatial distribution of the clouds.  They suggest that the HVCs are the unused building blocks of the Local Group, falling towards its barycenter.  Braun \\& Burton~(\\cite{braun99}, hereafter BB99)   reached  similar conclusions based on a study of a distinct subset of the HVC population.  By restricting their attention to compact, isolated CHVCs, they exclude the contribution of the nearby, less representative clouds.  The hypothesis is that the compact clouds might be the distant counterparts of the nearby, large angular size complexes. The compact sample also shows a natural preference for the Local Group Standard of Rest system, wherein its velocity dispersion ($88\\rm\\;km\\;s^{-1}$) is lower than in either the LSR or GSR frames. The CHVCs even allow definition of a new coordinate system in which a global minimum of the velocity dispersion ($69\\rm\\;km\\;s^{-1}$) is obtained. This system agrees with the Local Group system at about the one sigma level.  Furthermore, analysis of high resolution images of sixteen of the CHVCs provide several independent indications of distances of between 150 and 850 kpc (Braun \\& Burton~\\cite{braun00}; Burton et al. \\cite{burton01a}, \\cite{burton01b}).  The BB99 sample was obtained by visual inspection of the Leiden/Dwingeloo Survey (LDS) of the local \\hi sky carried out by Hartmann \\& Burton (\\cite{hartmann97}).  The LDS surveyed the sky as far south as the Dwingeloo horizon, that is to a declination of~$-30^\\circ$; lacking information on the more southern declinations, the BB99 conclusions were based on an incomplete sample.  A major improvement of the CHVC study would be an extension to the whole sky. Its high sensitivity and fully--Nyquist sampling makes the recently completed Parkes All--Sky Survey, HIPASS, (Barnes et~al. \\cite{barnes01}) ideal for extending the CHVC sample.  To create an all--sky resource which is as homogeneous as possible, an automated algorithm has been developed and is described here.   This paper also discusses  application of the algorithm to the LDS;  a separate paper (Putman et~al. \\cite{putman02}) gives the results from applying the algorithm to the HIPASS southern--hemisphere data.  Our discussion is organized as follows. We begin by describing the data used and the importance of obtaining confirming observations in \\S\\ref{sect:data}, proceed with a description of the algorithm and selection criteria in \\S\\ref{sect:algorithm}, present a catalog of both compact and extended high--velocity clouds in \\S\\ref{sect:results}, and conclude with a brief discussion of the global properties of the cataloged objects in \\S\\ref{sect:discussion}. ",
        "conclusions": "\\label{sect:discussion}  \\subsection{What defines a CHVC?}  The concept of a distinct class of compact, isolated high--velocity clouds has emerged from the visual inspection of large area images of good sensitivity and spatial sampling like the LDS in the north and the HIPASS in the south. There appears to be a class of high-contrast features which are at best only marginally resolved with half degree angular resolution and that can not easily be distinguished from the \\hi signature of an external galaxy. The \\hi signature of an external galaxy in the LDS (see Figure 15 of Hartmann \\& Burton \\cite{hartmann97}) is a moderately high peak column density of a few times 10$^{19}$ cm$^{-2}$ or more (averaged over the 36 arcmin beam), an \\hi FWHM linewidth varying between 20 and 375 \\kms commensurate with the galaxy mass and inclination, and a {\\it sharply bounded angular extent}, such that emission at column densities above a few times 10$^{18}$ cm$^{-2}$ is confined to less than about 2$^\\circ$ diameter.  This last fact is of particular physical relevance, since the precipitous decline in \\hi column density seen at the edges of nearby galaxies can be understood as arising from photo--ionization due to the intergalactic radiation field. Detailed studies of individual systems (NGC3198 by Maloney \\cite{maloney93}, M33 by Corbelli \\& Salpeter \\cite{corbelli93}) show that an exponential decline in neutral column density, with a scale--length of about 1~kpc sets in below a critical column density of about 2$\\times$10$^{19}$ cm$^{-2}$.  This neutral column density is comparable to the Warm Neutral Medium layer (with a temperature of about 8000~K) required to provide sufficient shielding from UV and soft X-ray radiation such that condensation of Cool Neutral Medium clumps (with temperature of about 100~K) can take place (e.g. Wolfire et al. \\cite{wolfire95a}) for ambient thermal pressures comparable to, or less than, those found in the solar neighbourhood. The picture that emerges is a nested structure of CNM cores, shielded by WNM cocoons and surrounded by a Warm Ionized Medium halo. For nearby galaxies the WNM cocoon at the edge of the gaseous disk has an observed exponential scale-length of about 1~kpc (Maloney \\cite{maloney93}; Corbelli \\& Salpeter \\cite{corbelli93}). The exponential scale-length of the partially ionized WNM cocoon is determined by both the total column density distribution and the incident ionizing spectrum. For smooth distributions of total column density and a wide range of power--law spectral indices, Corbelli \\& Salpeter (\\cite{corbelli93}) find neutral scale--lengths of about 1~kpc. The reason that this transition zone is so much more extended than in a classical \\hii region is found in the wide range of photon energies, and hence penetrating depths, of the ionizing spectrum.  It is not   obvious that all of these considerations need apply to the population of high--velocity clouds. If the HVCs are near the Galactic disk, then they will be subject to severe tidal distortions of their intrinsic gas distributions and be exposed to variable ionizing radiation levels depending on local circumstances. Due to changing physical conditions it is conceivable that thermal and pressure equilibrium may not be achieved. The requirements for achieving such equilibrium have been considered by Wolfire et al. (\\cite{wolfire95b}), who conclude that even for heights above the Galactic plane of only 3~kpc, thermal equilibrium should be marginally achieved within the WNM and easily within the CNM at an infall velocity as high as 150~\\kms. At larger distances the requirements that the thermal timescale be shorter than the time to experience significant pressure variations are even more easily satisfied. It thus seems reasonable to expect that high--velocity clouds might also be described by CNM clumps within a WNM shielding cocoon surrounded by WIM halos. If such objects occur at distances greater than about 60~kpc, tidal effects would be less disruptive, the UV radiation field would be more nearly isotropic, and an \\hi concentration of about 2~kpc diameter would have an angular size of less than 2$^\\circ$.  \\subsection{Application of a digital filter to the LDS}  Motivated by the distinctive visual appearance of sharply bounded \\hi peaks in the LDS data, we have considered whether a simple digital filtering might not serve to isolate a subset of these features from the general \\hi emission of the Galaxy and the HVC complexes, and in so doing offer independent confirmation of the results achieved with the procedure outlined in \\S\\ref{sect:algorithm}.   We considered a two--dimensional discrete derivative convolving kernel, consisting of a delta function at the origin together with a negative ring of unit integral at a radius of 1$^\\circ$. After evaluating the convolution of each channel map with this kernel, we normalize the result by the original intensity, wherever this exceeds some minimum significance: \\begin{equation} F(x,y) =  \\frac{1}{I} \\frac{dI}{dr} \\end{equation} for $I > n\\sigma$. As an illustration, this filter was applied to an Aitoff-projected cube of the LDS data centered at $(l,b)=(123\\deg,0\\deg)$ after a velocity smoothing to 25 \\kms~FWHM, and restricted to intensities exceeding 10$\\sigma$, corresponding to 0.14~K brightness. The peak filtered response was then determined along each spectrum; this is illustrated as the greyscale in Figure \\ref{fig:skydis} for values exceeding 0.85 in units of the normalized derivative. Overlaid on the peak filter values are square symbols at the positions of 59 cataloged external galaxies.  Of these 59 nearby \\hi--emitting galaxies, 57 are listed in Table 4 of Hartmann \\& Burton (\\cite{hartmann97}); the remaining two are Dwingeloo 1, discovered by Kraan--Korteweg et al. (\\cite{kraan94}), and Cepheus 1, discovered by Burton et al. (\\cite{burton99}).  Open circles in this figure show the positions of the 116 CHVCs which we catalog here; crosses show the positions of 72 CHVC candidates which were not confirmed in subsequent observations, and triangles show the positions of 84 CHVC candidates for which no confirming observations have yet been obtained.  \\begin{figure*} \\centering \\caption{Distribution on the sky of the compact anomalous--velocity objects (CHVCs) found by application of the search algorithm and selection criteria to the Leiden/Dwingeloo Survey. The background greyscale is the output of a spatial derivative filter to the LDS data (matched to the median cloud linewidth and angular size) where this exceeds 0.85 in units of the normalized derivative. Overlaid are square symbols at the positions of 59 cataloged external galaxies, open circles for the 116 CHVCs which we catalog here; crosses at the positions of 72 CHVC candidates which were not confirmed in subsequent observations, and triangles at the positions of 84 CHVC candidates for which no confirming observations have yet been obtained.  The CHVCs are not strongly clumped in the northern sky; and in particular show no concentration toward the extended HVC complexes or Magellanic Stream but do show some concentration toward the region near (l,b)=(25$\\deg$,$-$30$\\deg$) which has previously been called the Galactic Center Negative velocity population. The solid curve indicates the nominal $\\delta = -30^\\circ$ cut-off of the LDS coverage.  } \\label{fig:skydis} \\end{figure*}  It is striking how the vast majority of both normal Galactic and HVC complex emission has been eliminated by this simple filtering. About half of the peaks returned by the normalized derivative filter have an over--plotted symbol corresponding either to a cataloged external galaxy or to a CHVC. The 18 cataloged galaxies and 39 CHVCs which do not have a strong filter response can all be understood in terms of a poorly matched velocity width (since objects well--matched with a 25 \\kms~FWHM were selected in this case) or a very low peak brightness (that is, below the cut--off of 0.14~K). In addition to the filter peaks with over--plotted symbols there are a similar number, about 140 peaks, exceeding 0.85 in the normalized spatial derivative without an overplotted identification. Closer inspection of these additional filter peaks reveals that they all correspond to sub--structure within more extended HVC complexes. These are quite distinct from the CHVCs in that they are not are not sharply bounded at a column density of 1.5$\\times$10$^{18}$ cm$^{-2}$, but instead are connected (at least in projected position--velocity space) to more extended HVC complexes at a higher column density.  \\subsection{Brief remarks on the spatial and kinematic deployment of the northern CHVCs}  In a separate paper (De Heij et al. \\cite{deheij02}), we  analyse the LDS CHVC sample merged with the sample identified by Putman et al. (\\cite{putman02}) in the southern hemisphere HIPASS material.  Here we briefly note the global properties of the compact objects found in the LDS.  Figure \\ref{fig:skydis} shows the distribution on the sky of the compact, isolated objects cataloged in Table~\\ref{table:CHVC}. The objects are distributed rather uniformly across the northern sky. In particular, there is no clear sign of a preference in the sky distribution either for the Galactic disk, or for most of the known high--velocity cloud complexes, or the Magellanic Stream. One exception to this general conclusion is an apparent concentration centered near (l,b)=(25,$-$30), a region which Wakker \\& Van Woerden (\\cite{wakker91}) identify with their Galactic Center Negative Velocity Population of HVCs. This region appears to be particuarly rich in compact objects statisfying very stringent requirements for isolation in column density. For comparison, the distribution of all nearby galaxies detected in the LDS data is also plotted in the figure.  Figure~\\ref{fig:veldis} shows the velocities of all of the compact clouds cataloged in Table~\\ref{table:CHVC}, calculated for three different reference frames and plotted as function of the Galactic longitude.  The velocity dispersion of the sample decreases from 174 \\kms, to 101 \\kms, and then to 97 \\kms, in going from the reference frame of the Local Standard of Rest, to that of the Galactic Standard of Rest, and then to that of the Local Group Standard of Rest.  The average velocity changes from $-191$ \\kms, to $-127$ \\kms, and then to $-117$ \\kms, for the three different reference frames, respectively. Although the differences between the Galactic and Local Group system are not very large, the values are lower in the Local Group coordinate system. Inclusion of the data of the southern hemisphere in the analysis has to indicate if the small difference is significant or not.  The difference of 10 \\kms~between the velocity dispersion in the LGSR reference frame found by BB99 and that found here is probably due to the differences in the samples used, with the BB99 one being smaller.   \\begin{figure} \\caption{Kinematic deployment of the compact--object sample plotted as a function of the galactic longitude for three different kinematic reference frames, compared with the kinematic distribution of Local Group galaxies.  The CHVCs are shown as triangles; the less--constrained :HVC and ?HVC objects as diamonds and squares, respectively.  The kinematic distribution of the Local Group galaxies tabulated by Mateo (\\cite{mateo98}) is traced by the filled circles. There is a decrease in the average velocity of the compact--object ensemble, as well as in the velocity dispersion, when progressing from the Local Standard of Rest system (upper panel), to the Galactic Standard of Rest system (middle panel), to the Local Group Standard of Rest system (lower panel). } \\label{fig:veldis} \\end{figure}  \\subsection{Summary and Conclusions}  An automated procedure has been developed to extract anomalous--velocity clouds from \\hi survey material.  We have applied the algorithm to the Leiden/Dwingeloo survey and catalog the properties of a total of 917 HVC features with a deviation velocity in excess of 70~\\kms. Since the algorithm requires the existence of a local maximum in position--velocity which is distinct from the bulk of Galactic emission, the catalog can not be usefully extended to include the intermediate velocity clouds which are strongly blended with the Galaxy.  We have searched our HVC catalog for all isolated clouds, defined by having a lowest significant column density contour (3$\\sigma\\sim1.5\\times$10$^{18}$ cm$^{-2}$) which is (1) closed, with its greatest radial extent less than $10\\deg$ by $10\\deg$; and (2) not elongated in the direction of any nearby extended emission. A total of 116 objects have been tabulated which at least partially satisfy these criteria. Independent confirmation is available for all of these clouds, some of which appear in only single spectra in the LDS.  Of the 116 clouds, 54 had been identified as CHVCs by Braun \\& Burton (\\cite{braun99}), some others had been detected in earlier work referenced in the catalog but are confirmed as isolated by the data and analysis presented here, and others are reported here for the first time.  Although objects as large as $10\\deg$ were permitted by our selection criteria, the resulting distribution is strongly peaked at a median value of $1\\deg$ FWHM and has a maximum observed diameter of only $2.\\fdg2$. Isolated HVCs are observed to be relatively compact. Conversely, although the well known HVC complexes exhibit a wealth of small-scale structure that is comparable in angular scale ($\\sim1\\deg$) these structures are {\\it not isolated}. The local maxima within the HVC complexes are surrounded by extended emission with column densities in the range $5-20\\times$10$^{18}$ cm$^{-2}$.  The significance of high galactic latitude \\hi features which are isolated in column density down to a level as low as $1.5\\times$10$^{18}$ cm$^{-2}$, is that this is about an order of magnitude lower than the critical column density identified at the edges of nearby galaxies (Maloney \\cite{maloney93}, Corbelli \\& Salpeter \\cite{corbelli93}), $\\sim$2$\\times$10$^{19}$ cm$^{-2}$, where the ionized fraction is thought to increase dramatically due to the extragalactic radiation field. Unless very contrived geometries are invoked of some unseen population of high column density absorbers, these objects will need to provide their own shielding to ionizing radiation. (This point has also been made by Hoffman et al. \\cite{hoffman01}.) Self-consistent calculations of the ionization balance in the shielding layer exposed to a power--law extragalactic ionizing photon field yield a typical \\hi exponential scale--length of 1~kpc (Corbelli \\& Salpeter \\cite{corbelli93}).  The small median angular size of the CHVCs, of about $1^\\circ$ FWHM, might then imply substantial distances, greater than about 120~kpc.  The kinematic and spatial deployment of the enlarged sample shows that both the Galactic Standard of Rest and the Local Group Standard of Rest frames substantially lower the velocity dispersion of the population. This also lends support to the hypothesis that the CHVCs are a distant component of the high--velocity cloud phenomenon, at substantial distances and with a net in--fall towards the Galaxy or the Local Group barycenter.  A more complete analysis of the kinematic and spatial deployment of the all--sky CHVC ensemble incoporates both the LDS sample and the results of a comparable search by Putman et al. (\\cite{putman02}) in the HIPASS data, and is reported separately by De Heij et al. (\\cite{deheij02})."
    },
    "0201/astro-ph0201005_arXiv.txt": {
        "abstract": "Pulsars are rotating neutron stars, sweeping the emission regions from the magnetic poles across our line of sight. Isolated neutron stars lose angular momentum through dipole radiation and (possibly) particle winds, hence they slow down extremely steadily, making them amongst the most reliable timing sources available. However, it is well known that younger pulsars can suffer glitches, when they suddenly deviate from their stable rotation period. On 2000 January 16 (MJD 51559) the rate of pulsation from the Vela pulsar (B0833-45) showed such a fractional period change of {\\rm $3.1\\times 10^{-6}$}, the largest recorded for this pulsar. The glitch was detected and reported by the Hobart radio telescope. The speedy announcement allowed the X-ray telescope, Chandra, and others, to make Target of Opportunity observations. The data placed an upper limit of 40 seconds for the transition time from the original to the new period. Four relaxation timescales are found, which are believed to be due to the variable coupling between the crust and the interior fluid.  One is very short, about 60 seconds; the others have been previously reported and are 0.56, 3.33 and 19.1 days in length. ",
        "introduction": "Observations of pulsar glitches, in addition to providing insights into the phenomenon itself, offer one of the few probes of neutron star structure, and thus the physics of ultra-dense matter. Vela is the brightest known radio pulsar, and as it is at a declination of -45$^o$, it is above the horizon at the Hobart Radio Observatory (Mount Pleasant) for more than 18 hours a day.  It undergoes large glitches in pulse rate every few years and so provides an excellent probe of the physics of glitches.   Since 1981 the University of Tasmania has devoted a 14m diameter antenna at its Mt Pleasant Observatory to measurements of arrival times of pulses from the Vela pulsar. During this time we have observed 7 large glitches, or sudden decreases, in the period of the pulsar \\citep{vela_1983,vela_1987,christ_nat,ppp_vela}.  The telescope has a single pulse observing system whose speed and sensitivity have been enhanced in order to answer a number of questions; how quickly does the crust accelerate to the new period during a glitch, how soon does the recovery from the glitch start, and what is the form of this recovery?  On 2000 January 16 (MJD 51559) the rate of pulsation jumped with a fractional period change of {\\rm $3.1\\times 10^{-6}$}, the largest recorded for this pulsar. The glitch was automatically detected and we issued an IAU telegram \\citep{iau_7347} within 12 hours, allowing the X-ray telescope, Chandra, to make Target of Opportunity (TOO) observations \\citep{vela_xray_pwn}. These observations have so far failed to find the signature of neutron star heating, which was the driver for the TOO, but have produced spectacular images of the X-ray pulsar wind nebula. ",
        "conclusions": ""
    },
    "0201/astro-ph0201233_arXiv.txt": {
        "abstract": "The origin of {\\it early superhumps}, which are double-wave periodic modulations seen only during the earliest stage of WZ Sge-type outburst, has not been well understood.  Based on recent discovery of two-armed arch-like patterns on Doppler tomograms in conjunction with early superhumps, we propose a new interpretation on the origin of early superhumps, following the new interpretation by \\citet{sma01tidal} and \\citet{ogi01tidal} of the two-armed pattern seen in IP Peg.  If we consider irradiation of the elevated surface of the accretion disk caused by vertical tidal deformation, we can consistently explain the observed features on Doppler tomograms and photometric waves at the same time. We interpret that a combination of low mass-ratios ($q$) and low mass-transfer rates, necessary to give rise to these tidal effects, discriminate WZ Sge-type stars from other SU UMa-type dwarf novae. Based on recent fluid calculations, such an effect would be observable in higher $q$ systems.  We interpret that RZ Leo is an example of such objects. ",
        "introduction": "Dwarf novae are a class of cataclysmic variables (CVs), which are close binary systems consisting of a white dwarf and a red dwarf secondary transferring matter via the Roche-lobe overflow.  WZ Sge-type dwarf novae (cf. \\cite{bai79wzsge}; \\cite{dow81wzsge}; \\cite{pat81wzsge}; \\cite{dow90wxcet}; \\cite{odo91wzsge}; \\cite{kat01hvvir}) are a small subgroup of dwarf novae characterized by the long ($\\sim$ 10 yr) outburst recurrence time and the large ($\\sim$ 8 mag) outburst amplitude. They are a subclass of SU UMa-type dwarf novae (cf. \\cite{war95suuma}).  The most remarkable signature of WZ Sge-type outbursts is the presence of ``early superhumps\"\\footnote{ This feature is also referred to as {\\it orbital superhumps} \\citep{kat96alcom} or {\\it outburst orbital hump} \\citep{pat98egcnc}. } during their earliest stage of superoutbursts.  Early superhumps have a period extremely close to that of the binary period\\footnote{ The best-established case is the 2001 superoutburst of WZ Sge (\\cite{ish02wzsgeletter} and references therein).  In a few other systems (e.g. AL Com: \\cite{kat96alcom}, \\cite{pat96alcom}, \\cite{nog97alcom}; EG Cnc: \\cite{kat97egcnc}; \\cite{mat98egcnc}; \\cite{pat98egcnc}), the periods of early superhumps have been found to be in good agreement with their quiescent photometric periods (most likely representing orbital periods). }, and commonly show double-humped profile (figure \\ref{fig:ecomp}; see also \\citet{kat96alcom} and \\citet{kat98super} for detailed discussions), in contrast to ordinary superhumps of SU UMa-type dwarf novae [see \\cite{war95suuma} for basic properties of SU UMa-type dwarf novae]. Early superhumps are the most discriminative feature of WZ Sge-type outbursts, and have not been detected in other dwarf novae (see \\cite{kat01wxcet} and \\cite{kat01hvvir}). Although several models have been historically proposed [e.g. enhanced hot spot \\citep{pat81wzsge}, immature form of superhumps \\citep{kat96alcom}, jet or thickened disk \\citep{nog97alcom}], none of them has successfully explained all the observed features of early superhumps.  \\begin{figure} \\begin{center} \\FigureFile(80mm,110mm){ecomp.eps} \\end{center} \\caption{Comparison of profiles of early superhumps of WZ Sge-type stars. The data are taken from: WZ Sge \\citep{ish02wzsgeletter}; RZ Leo \\citep{ish01rzleo}; AL Com \\citep{kat96alcom}; HV Vir \\citep{kat01hvvir}; EG Cnc \\citep{mat98egcnc}.  The phase in WZ Sge corresponds to the orbital phase.  Since the exact binary phases are unknown for the rest of the objects, the phases were arbitrarily taken so that the main hump maxima best match that of WZ Sge.  All stars exhibit characteristic double (or sometimes triple) humps with unequal amplitudes. } \\label{fig:ecomp} \\end{figure} ",
        "conclusions": ""
    },
    "0201/astro-ph0201469_arXiv.txt": {
        "abstract": "The connection between Gamma-Ray Bursts (GRBs) and their afterglows is currently not well understood. Afterglow models of synchrotron emission generated by external shocks in the GRB fireball model predict emission detectable in the gamma-ray regime ($\\gax 25$ keV). In this paper, we present a temporal and spectral analysis of a subset of BATSE GRBs with smooth extended emission tails to search for signatures of the ``early high-energy afterglow'', i.e., afterglow emission that initially begins in the gamma-ray phase and subsequently evolves into X-Ray, uv, optical, and radio emission as the blast wave is decelerated by the ambient medium. From a sample of 40 GRBs we find that the temporal decays are best described with a power-law $\\sim t^{\\beta}$, rather than an exponential, with a mean index $\\langle \\beta \\rangle \\approx -2$. Spectral analysis shows that $\\sim 20\\%$ of these events are consistent with fast-cooling synchrotron emission for an adiabatic blast wave; three of which are consistent with the blast wave evolution of a jet, with $F_{\\nu} \\sim t^{-p}$. This behavior suggests that, in some cases, the emission may originate from a narrow jet, possibly consisting of ``nuggets'' whose angular size are less than $1 / \\Gamma$, where $\\Gamma$ is the bulk Lorentz factor. ",
        "introduction": "Afterglow emissions from Gamma-Ray Bursts (GRBs) in the X-ray, optical, and radio wavebands are in good agreement with afterglow models of relativistic fireballs (\\cite{wijers97}, \\cite{galama98}, \\cite{waxman97}, \\cite{vietri97}). The observed afterglow spectrum is well-described as synchrotron emission that arises from the interaction of the relativistic blast wave with bulk Lorentz factor $\\Gamma_{0} \\sim 10^2 - 10^3$ with the ambient medium (\\cite{meszaros97}, \\cite{galama98}). The often highly variable gamma-ray phase of the burst may reflect the physical behavior of the fireball progenitor through collisions internal to the flow, i.e., internal shocks (\\cite{sari97}, \\cite{kobayashi97}). On the other hand, Dermer and Mitman (1999) have suggested a blast wave with an inhomogenous external medium. Heinz and Begelman (1999) have suggested an inhomogeneous bullet-like jet outflow that encounters the interstellar medium.  The precise relationship between the observed GRB and the afterglow emission is not well understood. GRBs recorded by the {\\it BeppoSAX} satellite suggest that the X-ray afterglow emission may be delayed in time from the main GRB (e.g., GRB970228, \\cite{costa00}) or may begin during the GRB emission (e.g., GRB980519, \\cite{intzand99}). In the latter case, it is not clear if the X-ray afterglow is a separate underlying emission component or a continuation of the GRB itself. The internal-external shock model presents a scenario in which emission from internal and external shocks may overlap in time. If the internal shocks reflect the activity of the progenitor, then the onset of the afterglow may be separated from the prompt gamma-ray emission. Since the nature of the progenitor is not known, the effect of the ambient medium on the emission from the progenitor is highly problematic. The model therefore does not prohibit internal and external shock emissions from overlap, while in other cases the afterglow emission may be delayed with respect to the GRB (e.g., \\cite{sari99a}, \\cite{meszaros99}, \\cite{vietri00}).  The detection of optical emission simultaneous with the gamma-ray emission of GRB990123 (\\cite{akerlof99}) provided the first evidence for two distinct emission components in a GRB; here, the prompt optical emission is believed to originate from synchrotron emission in the production of the reverse shock generated when the ejecta encounters the external medium (\\cite{galama99}, \\cite{sari99a}, \\cite{meszaros99}). The gamma-ray spectrum of GRB990123 can not be extrapolated from the spectral flux from the simultaneous optical emission, indicating that the optical and gamma-ray emission originate from two separate mechanisms (\\cite{briggs99}, \\cite{galama99}). Evidence for overlapping shock emission was also found in GRB980923 (\\cite{giblin99a}), where a long power-law decay tail ($\\sim t^{-1.8}$) was observed in soft gamma-rays (25-300 keV). Two separate emission components are favored in this burst because the spectral characteristics of the tail were markedly different from those of the variable main GRB emission. The spectrum in the tail is consistent with that of a slow-cooling synchrotron spectrum, similar to the behavior of low-energy afterglows (e.g., \\cite{bloom98}, \\cite{vreeswijk99}).  The gamma-rays produced by internal shocks and the soft gamma-rays of the ``afterglow'' may therefore overlap, the latter having a signature of power-law decay in the synchrotron afterglow model. If this is the case, at least some GRBs in the BATSE database should show signatures of the early external shock emission. These events would contain a soft gamma-ray (or hard X-ray) tail component that decays as a power-law in their time histories, possibly superposed upon the variable gamma-ray emission. It has been shown that the peak frequency of the initial synchrotron emission, which depends on the parameters of the system (see $\\S 2$), can peak in hard X-rays or gamma-rays (\\cite{meszaros92}). Further, it may be possible to see a smoothly decaying GRB that is the result of an external shock, i.e., the GRB itself is a ``high-energy'' afterglow. For such GRBs, the subsequent afterglow emission in X-rays and optical would then simply be the evolution of the burst spectrum. A situation like this might arise when the progenitor generates only a single energy release (i.e., no internal shocks).  It is well known that the temporal structures of GRBs are very diverse and often contain complex, rapid variability. However, some bursts exhibit smooth decay features that persist on timescales as long as, or even longer than, the variable emission of the burst. Our investigation focuses on the combined temporal and spectral behavior of a sample of 40 BATSE GRBs that exhibit smooth decays during the later phase of their time histories. Many of these events fall into a category of bursts traditionally referred to as ``FREDs'' (Fast Rise, Exponential-like Decay), bursts with rapid rise times and a smooth extended decay (\\cite{kouveliotou92}). In $\\S 2$ we present temporal and spectral properties of the afterglow synchrotron spectrum. In $\\S 3$ we examine the temporal behavior and spectral characterisitics of the decay emission for the events in our sample and compare their spectra with the model synchrotron spectrum. A color-color diagram (CCD) technique is also applied to systematically explore the spectral evolution of each event. In $\\S 4$ we present a set of high-energy afterglow candidates, followed by a discussion of our results in the framework of current fireball models. ",
        "conclusions": "Our temporal and spectral analysis of the smooth extended gamma-ray decay emission in GRBs has shown evidence of signatures for early high-energy afterglow emission in gamma-ray bursts. The extended decay emission is best described with a power-law function $F_{\\nu} \\propto t^{\\beta}$ rather than an exponential, similar to the results of Ryde and Svensson (2001) who studied the decay phase of a sample of GRB pulses with a broad range of durations. From our sample of 40 events, we find $\\langle \\beta \\rangle \\approx -2$ for long, smooth decays. Color-color diagrams have provided a qualitative interpretation of the burst spectral evolution and allow a simple comparison with the evolution expected from the synchrotron model as well as comparison of spectral evolution among GRBs. The CCD patterns and the spectral analysis indicate that $\\sim 20\\%$ of the events in our sample are consistent with synchrotron emission expected from an external shock. Interestingly, three of these events have decay rates consistent with that expected from the evolution of a jet, $F_{\\nu} \\sim t^{-p}$. Because the break is essentially at the onset of deceleration, the jet must, at least, be very narrow, since $\\theta_{c} < 1/\\Gamma$. Table 4 suggests that in some cases the break occurs at a later time, so that the prompt emission we observe is pre-break, $\\theta_{c} > 1 / \\Gamma$ and consistent with spherical geometry. A possible scenario is one in which the ejecta is very grainy, where the nuggets in the ejecta are smaller than $1/\\Gamma$, similar to the model discussed by Heinz and Begelman (1999). Huang et al.\\ (1999) (see also Wei and Lu 2000) have shown that the break in the light curve is more of a smooth transition due to the off-axis emission of a jet with no angular dependence. The steep light curve can only occur if the angular size of the nugget is less than $1/\\Gamma$.  Connaughton (2001) has investigated the average late-time temporal properties of GRBs observed with BATSE and found statistically significant late time power-law decay emission that softens relative to the initial burst emission, suggesting the existence of early high-energy afterglow. Other studies using PHEBUS (\\cite{tkachenko00}) and APEX (\\cite{litvine00}) bursts show similar behaviors in late-time GRB light curves. Collectively, these studies strongly suggest that the afterglow emission may overlap or be connected to the prompt, variable burst emission. On the other hand, it is clear that not all GRBs exhibit such behavior. In some cases, the initial gamma-ray flux from the external shock may simply be too low to detect (e.g., see Figure 4 in Giblin et al.\\ 2000). In other cases, the bulk Lorentz factor may be too low to generate the gamma-ray photons upon impact with the surrounding medium.  As the number of afterglow/counterpart detections increases, the relationship of the afterglow emission to the gamma-rays released in the initial phase of the burst can be studied systematically. The capabilities of {\\it Swift} (\\cite{gehrels02}) will allow broad spectral coverage using three co-aligned instruments (BAT, XRT, and UVOT) during the gamma-ray phase and early afterglow phase of the burst and facilitate the distinction between the GRB and the onset of the afterglow based on temporal {\\it and} spectral information. With well-constrained spectral and temporal parameters in hand, plots of temporal index vs.\\ spectral index can be readily constructed and thus provide information on the geometry of the fireball and definitively test the internal/external shock model for GRBs."
    },
    "0201/astro-ph0201143_arXiv.txt": {
        "abstract": "We present the 158 standard stars that define the $u'g'r'i'z'$ photometric system.  These stars form the basis for the photometric calibration of the Sloan Digital Sky Survey (SDSS).  The defining instrument system and filters, the observing process, the reduction techniques, and the software used to create the stellar network are all described.  We briefly discuss the history of the star selection process, the derivation of a set of transformation equations for the $UBVR_{\\rm c}I_{\\rm c}$ system, and plans for future work. ",
        "introduction": "We present the newly established standard star network for the $u'g'r'i'z'$ filter system \\citep[see][]{fuk96}.  This standard star network was developed at the U.S.  Naval Observatory, Flagstaff Station. These stars form the basis for the photometric calibration of the Sloan Digital Sky Survey (SDSS).  The SDSS uses a 2.5-m telescope at Apache Point Observatory (APO) to produce a five-band, photometrically calibrated digital imaging survey of $\\pi$ steradians (10,000 square degrees) of the Northern Galactic Cap \\citep{GCRS98,york00} as one of its major data products.  It is not our purpose here to describe in detail the full end-to-end process of calibrating the SDSS photometric data.  That is the topic of a future paper.  Here, we merely wish to present a self-contained description of the standard star network upon which the SDSS photometry is based.  We do note, however, that one of the targets of the SDSS is to achieve a level of photometric uniformity and accuracy such that the system-wide rms errors in the final SDSS photometric catalog will be less than 0.02~mag in $r'$, 0.02~mag in $(r'-i')$ and $(g'-r')$, and 0.03~mag in $(u'-g')$ and $(i'-z')$, for objects bluer than an M0 dwarf.  To meet this target, internal goals were set for the accuracy of the primary standard star system: the uncertainty in the mean calibrated magnitudes for any given primary standard star should be $\\le$1.5\\% at $u'$, $\\le$ 1\\% in $g', r'$ and $i'$, and $\\le$1.5\\% at $z'$.  As we will show later in this paper, we more than meet these goals for all but a handful of stars.  In addition, we must mention that, due to small-but-significant differences between the USNO and 2.5-m filters, the final 2.5-m SDSS published photometry will likely differ systematically from the $u'g'r'i'z'$ system at the few percent level for $g'r'i'$ and slightly worse for $u'$ and $z'$ \\citep[see][]{sto01}.  When transformation equations between the $u'g'r'i'z'$ system and the 2.5-m SDSS photometry have been robustly determined, they will be promptly made available to the astronomical community.  (Note: the intended accuracy of these transformation equations is included within the above-mentioned error budget for the photometric calibrations of the final SDSS imaging catalog.)  The nomenclature used in this system differs slightly from the traditional photometric literature.  This was done to avoid confusion with existing SDSS papers and nomenclature.  In the traditional sense, Vega ($\\alpha$ Lyr) is the ultimate ``fundamental'' standard.  However, in this paper we refer to three subdwarf stars which were used to set the initial system zeropoint as ``fundamental'' with the other 155 stars of the system being referred to as ``primary'' stars.  The term ``secondary'' is used within the SDSS nomenclature to refer to the photometric system transfer patches --- pieces of the sky that are observed by a 0.5-m telescope that are used to transfer the photometric solution to the main survey imaging telescope.  In the following sections we present details of the standard star development program.  We describe the instrumentation and filter system in \\S2, selection of the initial set of stars in \\S3, and a brief overview of the reduction software that was used to develop the network in \\S4.  Final results for the inital set of $u'g'r'i'z'$ primary standard stars are presented in \\S5, and we discuss future extensions to this system in \\S6. ",
        "conclusions": ""
    },
    "0201/astro-ph0201375_arXiv.txt": {
        "abstract": "The Sunyaev Zel'dovich effect (SZ effect) is a complete probe of ionized baryons, the majority of which are likely hiding in the intergalactic medium.  We ran a $512^3$ $\\Lambda$CDM simulation using a moving mesh hydro code to compute the statistics of the thermal and kinetic SZ effect such as the power spectra and measures of non-Gaussianity. The thermal SZ power spectrum has a very broad peak at multipole $l\\sim 2000-10^4$ with temperature fluctuations $\\Delta T \\sim 15\\mu$K. The power spectrum is consistent with available observations and suggests a high $\\sigma_8\\simeq 1.0$ and a possible role of non-gravitational heating. The non-Gaussianity is significant and  increases the cosmic variance of the power spectrum by a factor of $\\sim 5$ for $l<6000$.  We explore optimal driftscan survey strategies for the AMIBA CMB interferometer and their dependence on cosmology.  For SZ power spectrum estimation, we find that the optimal sky coverage for a $1000$ hours of integration time is several hundred square degrees. One achieves an accuracy better than $40\\%$ in the SZ measurement of power spectrum and an accuracy better than $20\\%$ in the cross correlation with Sloan galaxies for $2000<l<5000$.   For cluster searches, the optimal scan rate is around $280$ hours per square degree with a cluster detection rate 1 every $7$ hours, allowing for a false positive rate of $20\\%$ and better than $30\\%$ accuracy in the cluster SZ distribution function  measurement. ",
        "introduction": "Big bang nucleosynthesis (BBN) and cosmic microwave background (CMB) experiments such as  Boomerang \\citep{Netterfield01} and DASI \\citep{Pryke01} predict that the ordinary baryonic matter accounts for about $5\\%$ of the total matter in the universe. But the directly observed components, such as stars, interstellar medium and intracluster gas, only account for  a few percent of this baryon budget, while more than $90\\%$ baryons have escaped direct detections \\citep{Persic92,Fukugita98}. This is known as the missing baryon problem.  The missing baryons are believed to be in the form of the intergalactic medium (IGM) and have been difficult to detect directly. To understand their state, such as density, temperature, peculiar velocity and metalicity, stands as a major challenge to both observation and theory, and is crucial to understand the thermal history of the universe and galaxy formation.  The IGM has various direct observable tracers. (1) Neutral hydrogen absorbs background light of quasars and produces the Lyman-$\\alpha$ forest. (2) Ionized electrons have thermal and peculiar motions and are capable of scattering CMB photons and generate secondary CMB temperature fluctuations, which are known as the thermal and kinetic Sunyaev Zel'dovich effects (SZ effect), respectively. Their precision measurements are becoming routinely available with the devoted CMB experiments such as AMIBA \\citep{AMIBA} and SZA (Sunyaev Zel'dovich array).  (3) Ionized electrons and protons interact with each other and emit X-rays through thermal bremsstrahlung and contribute to the soft X-ray ($0.5-2$ keV) background (XRB).  Several other tracers have been proposed, including X-ray absorption techniques \\citep{Perna98}, but they depend strongly on chemical compositions, and would be difficult to associate with direct IGM properties. These tracers probe different IGM phases and help to extract the IGM state. For example, the XRB flux upper limits place constrain on the gas clumpyness and suggest a potentially strong role of feedback \\citep{Pen99}.  Among these tracers, the SZ effect is a particularly powerful IGM probe. (1) It provides a complete sample of the intergalactic medium. All free electrons participate in Thomson scattering and contribute to the SZ effect. The Thomson optical depth from the epoch of reionization $z\\sim 10$ to the present is $\\tau \\sim 0.1$, which means that about $10\\%$ of CMB photons have been scattered by electrons. Since the number of CMB photons is much larger than the baryon number and the ionization fraction of our universe is larger than $99\\%$, most baryons contribute to the SZ effect. Compton scattering does not depend on redshift and is not affected by distance or the expansion of the universe. So, the SZ effect can probe the distant universe. Furthermore, the SZ effect does not strongly depend on gas density and probes a large dynamic range of baryon fluctuations. In contrast, the X-ray emission measure depends on density squared, and primarily probes the densest IGM regions at low redshift.  The Lyman-$\\alpha$ forest probes the neutral IGM, which in ionization equilibrium depends also on the square of the baryon density.  The neutral fractions only accounts for a tiny fraction of total baryons. Extrapolations over many orders of magnitude are required to understand the bulk of the baryons.  (2) It is straightforward to understand.  The simple dependence of the SZ effect on the ionized gas pressure or momentum does not put as stringent requirement on simulation resolution, nor as accurate an understanding of the gas state such as metalicity, temperature, velocity and ionization equilibrium, as X-rays and Lyman-$\\alpha$ properties do.  Pressure is the total thermal energy content of the gas, and finite volume flux conservative simulations are particular amenable to its modelling.  It is also easier to model analytically. (3) It has strong observational potential.  In this paper we will primarily consider the thermal SZ effect, which is easier to observe.  Its unique dependence on  frequency allows us in principle to disentangle the SZ effect from the contamination of the primary CMB and other noise sources \\citep{Cooray00b}. Hereafter, if not otherwise specified, the SZ effect always means the thermal SZ effect.  In SZ observations, all redshifts are entangled and smear some key information of the intervening IGM: its spatial distribution and time evolution. These properties around redshift $z\\sim 1$ are sensitive to many cosmological parameters and thus potentially promising to break degeneracies in cosmologies from the primary CMB experiments. One can hope to resolve the equation of the state of the dark energy. We are currently performing simulations for different cosmologies, which are degenerate in the primary CMB, to test the potential of such a procedure \\citep{Zhang02}. If combined with other observations, the SZ effect will become even more powerful. For example, with the aid of a galaxy photometric redshift survey, the evolution of the 3D gas pressure power spectrum and its cross correlation with the galaxy distribution can be extracted \\citep{Zhang01a}.  The above analysis depends on a detailed quantitative theoretical understanding of the SZ statistics.  Analytical approaches considered in the past strongly depend on various assumptions.  The Press-Schechter approach as adopted by \\citet{Cole88,Makino93,Atrio99,Komatsu99,Cooray00,Molnar00} requires ad hoc models for the gas profile in halos, whose shape and evolution are uncertain. The hierarchical method as proposed by \\citet{Zhang01a} strongly depends on the gas-dark matter correlation, which is also poorly understood. These estimates can be significantly improved by high-resolution simulations.  But past simulation results \\citep{Scaramella93,daSilva00,Refregier00, Seljak01,Springel01} disagreed on both amplitude and shape of the SZ power spectrum (see \\citet{Springel01} for a recent discussion). Simulation resolution may play a key role in this discrepancy since the SZ effect is sensitive to small structures, as discussed by \\citet{Seljak01}.  Differences between code algorithms may also cause part of the discrepancy. To address these problems, we ran the largest adaptive SZ simulation to date, a direct $512^3$ moving mesh hydro (MMH) code \\citep{Pen98a} simulation,  to reinvestigate the SZ statistics. For a better understanding of numerical and code issues in the SZ effect, our group is carrying out a series of simulations with identical cosmological parameters and initial condition but using different codes with different resolutions \\citep{Codes}.  Apart from simulation issues, a problem that has received little attention is analysis strategies for SZ data.  Several interferometric arrays such as AMIBA and SZA are under construction, but detailed data analysis models are still in their infancy.  With our simulated SZ maps, we can estimate the sensitivity and accuracy of SZ observations.  Given a sky scan rate, how accurately can the SZ power spectrum be measured? How many cluster can be found?  How accurately can the SZ decrement be determined for individual clusters? What are optimal strategies for these measurements?  In this paper, we will take AMIBA as our target to address these questions.  This paper is organized as follows: in \\S \\ref{sec:application}, we describe the SZ effect, its statistics and our method to analyze these statistics.  We then present our simulation results of these statistics and possible constrain from observations (\\S\\ref{sec:simulation}). In \\S\\ref{sec:AMIBA} we simulate AMIBA driftscan observations to estimate the accuracy of AMIBA measurement of these statistics.  We conclude in \\S \\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  We have performed the largest high resolution SZ simulations to date, and analysed the results.  We found further increases in the power spectrum relative to previous simulations, and significantly more small scale structures.  This trend is confirmed in Press-Schechter estimates, and suggests nominally increasing power spectra.  The small structure  behavior may not be a robust prediction, since non-gravitational effects will significantly modify those scales. We have examined the skewness and kurtosis on the sky maps, found a strong non-Gaussianity on subdegree scales and confirmed the log-normal distribution found in previous studies.  The Gaussian estimates of power spectrum sample variances are less severely affected, due to the averaging of many patches by each Fourier modes. But its effect on the power spectrum error analysis is significant.  We simulated SZ observations with AMIBA, and analyzed the sensitivity for different scan rates.  In a $1000$ hour survey, the optimal strategy for power spectrum estimations is to scan several hundred square degrees.  The SZ angular power spectrum measured in such survey can be determined to an accuracy of $\\sim 40\\%$ over a range of $2000 \\lesssim l \\lesssim 5000$.  A cross correlation with SDSS should allow an accuracy of $20\\%$ in the cross correlation measurement, which suggests that the time resolved measurement of the pressure power spectrum is then possible.  This scan rate results in a low  detection rate of clusters of galaxies, approximately one per $60$ hours with a false positive rate of 20\\%.  For the purpose of cluster search, the optimal scan rate is around $280$ hours per square degree, which could find $1$ cluster every $7$ hours, with the accuracy of $30\\%$ in the cluster $y$ measurement and $30\\%$ in the cluster CDF up to $N(y>y_p)\\sim 60\\  {\\rm deg}^{-2}$.  The predicted SZ power spectrum is consistent with recent indications from the CBI experiment \\citep{Mason01a,Sievers01} and the BIMA upper limit ($95\\%$ confidence). But it is higher than the BIMA $1$-$\\sigma$ result.  This may be a first indication of IGM non-gravitational feedback. Future blank sky surveys with data analysis  considering actual SZ non-Gaussianity will provide us with a quantitative understanding of the thermal history of the universe."
    },
    "0201/astro-ph0201361_arXiv.txt": {
        "abstract": "We present the first results of an ongoing project to study the morphological, kinematical, dynamical, and chemical properties of satellite galaxies of external giant spiral galaxies. The sample of objects has been selected from the catalogue by Zaritsky et al. (1997). The paper analyzes the morphology and structural parameters of a subsample of 60 such objects. The satellites span a great variety of morphologies and surface brightness profiles. About two thirds of the sample are spirals and irregulars, the remaining third being early-types. Some cases showing interaction between pairs of satellites are presented and briefly discussed. ",
        "introduction": "The standard model for the formation of large structures in the Universe predicts a hierarchical scenario in which the first generation of objects created correspond to subgalactic masses (Frenk et al. 1988).  The aggregation of such objects by accretion, merging, etc., over cosmological timescales would generate the wide variety of structures observed in the Universe. Mergers between two objects of similar masses destroy  disks (Barnes \\& Hernquist 1991) and seem to be one of the mechanisms that create ellipticals from spiral galaxies.  A merger between a giant galaxy and a small satellite tends to heat and thicken the disk (Vel\\'azquez \\& White 1999) and  could contribute to the growth of the bulge and an increase in the S\\'ersic $n$ parameter (S\\'ersic 1968) moving the galaxy toward earlier types (Walker, Mihos, \\& Herndquist (1996); Aguerri, Balcells, \\& Peletier 2001). This accretion of satellites can also trigger starbursts  (Mihos \\& Hernquist 1994), and produce counterrotating disks (e.g., Thakar \\& Ryden 1996). However, there is some controversy over whether or not the observed structure in disk galaxies is compatible with the number of such mergers expected in the standard model (Toth \\& Ostriker (1992); Zaritsky (1995); Weinberg (1997); Velazquez \\& White 1999).  The Milky Way and M31 are examples of bright spiral galaxies surrounded by several satellites. Semi-analytic models (Kauffmann, White, \\& Guideroni 1993) and  numerical simulation (Klypin et al. (1999); Moore et al.  1999) of hierarchical clustering galaxy formation predict a number of satellites an order of magnitude larger than observed in the Local Group. Several alternatives have been proposed to solve this discrepancy. It has been suggested (Klypin et al. (1999), Blitz et al. 1999) that these missing satellites were in the form of the compact high velocity clouds identified by Braun \\& Burton (1999). Bode et al. (2001) have analyzed a warm dark matter model in which the dark matter particles have non-zero velocities; these tend to suppress the structures on small scales. In this model  small halos should be formed by fragmentation instead of by aggregation. Using numerical simulations, Bode et al. show a reasonable agreement between the predictions of their model and observations on the number of satellites, the spatial distribution, and the epoch of formation. A possible reionization of the Universe at high redshift could also suppress the formation of small galaxies after that epoch (Bullock, Kravtsov, \\&  Weinberg (2000); Somerville (2001); Benson et al. 2001).  The existence of satellite galaxies in external systems has been known for a long time ($e.\\,g.$ Page (1952), Holmberg 1969). The most complete compilation and study of such objects was conducted by Zaritsky et al. (1997) who presented a catalogue containing 115 satellites orbiting 69 primary isolated spiral galaxies. Zaritsky et al. defined satellites as objects at projected distances $\\le 500$ kpc from their primaries, with differences in recessional velocity $\\le 500$ km s$^{-1}$ and  at least 2.2 mag fainter than their parent galaxy. The number of catalogued satellites per system is very small (typically 1--2) and for construction of the catalogue they should correspond to the brightest part of the luminosity function of these objects (see Pritchet \\& van den Bergh 1999). Although the completeness of this catalogue is hardly understood (see Zaritsky et al. 1993), it can be estimated to absolute limiting magnitude M$_B\\sim -15.5$. Zaritsky et al., have used the catalogue to demonstrate that the halos of the primaries extend to at least 200 kpc and have masses larger than $2\\times 10^{12}\\ M_\\odot$. Other observational studies have been conducted by Carignan et al. (1997), who have discovered and analyzed eight dwarf galaxies orbiting the massive lenticular galaxy NGC~5084. This group shows a nice example of one of the predictions of the standard models: i.e., there should be a net excess of satellites in retrograde orbits because, as the numerical simulations show (Quinn \\& Goodman 1986), it is more effective for the parent galaxy to accrete the satellites in direct orbits. Cote et al. (1997) have discovered 16 and 20 dwarf galaxies in the Sculptor and the Centaurus group respectively. Some controversy exist about the reality of the so-called ``Holmberg effect'', i.e., the fact that satellites tend to avoid the plane of the primary; while Zaritsky et al., claimed  the existence of the effect in their sample, Carignan et al., have not found it in their study of the NGC~5084 group.  In the study of these external satellites it is interesting to compare their properties with those observed in the Local Group; this will help us to understand the theoretical models for the formation of structures in galactic halos. We are conducting an ongoing observational  program which comprises broad band photometry in the optical ($B$, $V$, $R$, and $I$) and infrared  ($J$ and $K'$), and in the H$\\alpha$ narrow band for both the parent and the satellite galaxies, taken from the compilation by Zaritsky et al. (1997). The aim of this study is to determine the structure,  dynamics, star formation rates, and stellar populations to constrain the age and metallicity of such systems.  Here, we concentrate on the morphology and structural parameters of the satellite galaxies. ",
        "conclusions": "The results presented here show that satellite galaxies present a wide variety of morphological types, sizes, and brightness. The Hubble types go from pure ellipticals to irregulars. About 35 objects have been classified as spiral galaxies.  The largest group in the catalogue is NGC 1961 and its five satellites. All these satellites are luminous ($-19.6\\le M_B \\le -18.1$) spiral galaxies. This group is notably different from the Local Group in which only M33 and LMC have  magnitude in this range. The other large groups in the sample are those formed by NGC 4030 and NGC 4541, which have four detected satellites each. This last is a notable group in which two of the galaxies show clear interaction (see below), and another, NGC 4541d, shows strong asymmetries. There are two clear cases of interaction between pairs of satellites: NGC 2718a and NGC 2718b, and NGC 4541b and NGC 4541e, respectively. This is illustrated in Figure 2. In the first case the interaction shows up in the form of a bridge connecting both galaxies. In the case of the second interacting pair, two tidal tails seem to emerge from NGC 4541b. One of them points directly to NGC 4541e and  it seems clear that its origin is the interaction with this galaxy; the other is nearly in the opposite direction and could reveal stripping from a previous orbit. The asymmetries shown by NGC 4541e could be also related with this interaction. Another notable case is NGC 4030c classified as Irr, in which several extended objects are surrounded by a common diffuse structure resembling a case of strong interaction or merger.  In a forthcoming paper, we will analyze the relationship between the different structural parameters, and between these parameters and the relative position and orientation of the satellites with respect to the parent galaxy. Considering the possible bias, completeness, and sampling of our data set, we will also conduct a detailed comparison between the properties found in our systems and those  in the Local Group."
    },
    "0201/astro-ph0201157_arXiv.txt": {
        "abstract": "MHD turbulence plays a central role in the physics of star-forming molecular clouds and the interstellar medium. {\\revised MHD turbulence in molecular clouds must be driven to account for the observed supersonic motions in the clouds, as even strongly magnetized turbulence decays quickly.  Driven MHD turbulence can globally support gravitationally unstable regions, but local collapse inevitably occurs.} Differences in the strength of driving and the gas density may explain the very different rates of star formation observed in different galaxies.  {\\revised Two types of comparisons to observations are reviewed.  First, the use of wavelet transform methods suggest that the driving comes from scales larger than observed molecular clouds.  Second, comparison of simulated spectral cubes from models to real observations suggests that Larson's mass-size relationship is an observational artifact.} The driving mechanism for the turbulence is likely a combination of field supernovae in star-forming sections of galactic disks, and magnetorotational instabilities in outer disks and low surface brightness galaxies.  Supernova-driven turbulence has a broad range of pressures with a roughly log-normal distribution.  High-pressure, cold regions can be formed even in the absence of self-gravity. ",
        "introduction": "{\\revised One of the big questions} in star formation is what determines the rate of star formation in galaxies?  Another, more pointed way of phrasing this question is to ask why the star formation rate in normal galaxies is so low, and why it varies so strongly, over orders of magnitude from low surface brightness galaxies, through normal galaxies, to starburst galaxies.  The free-fall time for gas at typical interstellar densities is \\begin{equation} t_{\\rm ff} = \\sqrt{\\frac{3\\pi}{32 G \\bar\\rho}} \\approx (3.4 \\times 10^7 \\mbox{ yr}) \\left(\\frac{n}{1 \\mbox{ cm}^{-3}}\\right)^{-1/2}, \\label{equ:tff} \\end{equation} where $\\bar\\rho$ is the mean mass density of the gas, $G$ the gravitational constant and $n=\\bar\\rho/\\mu$ the number density, with $\\mu=2.36 m_H$. Yet galactic ages range up to $10^{10}$~yr, and star formation continues today.  What has delayed star formation sufficiently to allow it to continue?  In what might be called the standard theory of star formation, magnetic fields are invoked to answer both of these questions.  If fields are strong enough, they can magnetostatically support clouds against collapse.  The star formation rate would then be determined by the rate of ambipolar drift of neutral gas past ions tied to the magnetic field towards the centers of self-gravitating cores \\cite{m77,s77}.  Furthermore, if the fields are strong enough that the Alfv\\'en speed $v_A$ reaches the rms velocity $v$, then strong shocks will be converted to MHD waves.  As linear Alfv\\'en waves are lossless, it was thought that motions remaining from the initial formation of the clouds might be enough to explain the observation of strongly supersonic motions in molecular clouds \\cite{am75}.  In this review I will explain why both of these ideas now appear questionable. ",
        "conclusions": "\\begin{itemize}  \\item Even relatively strong magnetic fields, with the field in equipartition with the kinetic energy, cannot prevent the decay of turbulent motions on dynamical timescales far shorter than the observed lifetimes of molecular clouds.  The significant kinetic energy observed in molecular cloud gas must be supplied more or less continuously.  \\item Supersonic turbulence strong enough to globally support a molecular cloud against collapse will usually cause {\\em local} collapse.  The turbulence establishes a complex network of interacting shocks.  The local density enhancements in fluctuations created by converging shock flows can be large enough to become gravitationally unstable and collapse. The probability for this to happen, the efficiency of the process, and the rate of continuing accretion onto collapsed cores are strongly dependent on the driving wave length and on the rms velocity of the turbulent flow, and thus on the driving mechanism.  \\item Interstellar clouds driven on large scales or without even global turbulent support very rapidly form stars in clusters.  On the contrary, in gas that is supported by turbulence, local collapse occurs sporadically over a large time interval, forming isolated stars.  The total star formation efficiency before the cloud dissolves due to stellar feedback or external shocks will probably be low. Thus, the strength and nature of the turbulence may be fully sufficient to explain the difference between the observed isolated and clustered modes of star formation.  \\item Magnetorotational instabilities may provide a base value for the velocity dispersion below which no galaxy will fall.  If that is sufficient to prevent collapse, little or no star formation will occur, producing something like a low surface brightness galaxy with large amounts of H~{\\sc i} and few stars.  In star-forming galaxies, however, clustered and field supernova explosions, predominantly from B~stars no longer associated with their parent gas, appear likely to dominate the driving, raising the velocity dispersion to some 10--15 km~s$^{-1}$.  \\item In a supernova-driven interstellar medium, we find a broad range of pressures with a log-normal distribution, and a substantial fraction of associated densities far from the thermal equilibrium values.  This limits the predictive usefulness of phase diagrams based on thermal equilibrium, although thermal equilibrium at the local pressure will still be the mildly favored state.  Gas pressures appear to be determined dynamically, and each individual parcel of gas seeks local thermal equilibrium at the pressure imposed on it by the turbulent flow.  Inferences that molecular clouds must be gravitationally bound because of their high observed confinement pressures are called into question by these results.  Regions with densities approaching the overall densities of GMCs, and pressures an order of magnitude above the average interstellar pressure appear in our simulations even in the absence of self-gravity.  \\end{itemize}   \\vspace{0.2in} I thank {\\revised the referee of this review, E. V\\'azquez-Semadeni, for a detailed and thoughtful report,} my collaborators M. A. de Avillez, J. Ballesteros-Paredes, D. Balsara, A. Burkert, F. Heitsch, J. Kim, R. S. Klessen, V. Ossenkopf, and M. D. Smith for their participation in different parts of the work reviewed here, and the organizers {\\revised of the conference} for their partial support of my attendance. This work was also partially supported by the NSF under CAREER grant AST99-85392 and by the NASA Astrophysical Theory Program under grant NAG5-10103. This research has made use of NASA's Astrophysics Data System Abstract Service."
    },
    "0201/astro-ph0201227_arXiv.txt": {
        "abstract": "Using BATSE and {\\em RXTE} observations from 1991 April to 2001 August we have detected 71 outbursts from 82 periastron passages of EXO 2030+375, a 42-second transient X-ray pulsar with a Be star companion, including several outbursts from 1993 August to 1996 April when the source was previously believed to be quiescent. Combining BATSE, {\\em RXTE}, and {\\em EXOSAT} data we have derived an improved orbital solution. Applying this solution results in a smooth profile for the spin-up rate during the giant outburst and results in evidence for a correlation between the spin-up rate and observed flux in the brighter BATSE outbursts. Infrared and H$\\alpha$ measurements show a decline in the density of the circumstellar disk around the Be star. This decline is followed by a sudden drop in the X-ray flux and a turn-over from a spin-up trend to spin-down in the frequency history. This is the first Be/X-ray binary which shows an extended interval, about 2.5 years, where the global trend is spin-down, but the outbursts continue. In 1995 the orbital phase of EXO 2030+375's outbursts shifted from peaking about 6 days after periastron to peaking before periastron. The outburst phase slowly recovered to peaking at about 2.5 days after periastron. We interpret this shift in orbital phase followed by a slow recovery as evidence for a global one-armed oscillation propagating in the Be disk. This is further supported by changes in the shape of the H$\\alpha$ profile which are commonly believed to be produced by a reconfiguration of the Be disk. The truncated viscous decretion disk model provides an explanation for the long series of normal outbursts and the evidence for an accretion disk in the brighter normal outbursts. Long-term multi-wavelength observations such as these clearly add considerably to our knowledge of Be/X-ray binaries and the relationship between optical, infrared and X-ray observations. ",
        "introduction": "Be/X-ray binaries are the most common type of accreting X-ray pulsar systems. They consist of a pulsar and a Be (or Oe) star, a main sequence star of spectral type B (or O) that shows Balmer emission lines (See e.g., Slettebak 1988 and Apparao 1994 for reviews.) The line emission is believed to be associated with an equatorial outflow of material expelled from the rapidly rotating Be star that probably forms a quasi-Keplerian disk near the Be star \\citep{Quirrenbach97,Hanuschik96}. X-ray outbursts are produced when the pulsar interacts with this disk. Be/X-ray binaries typically show two types of outburst behavior: (a) giant outbursts (or type II), characterized by high luminosities ($L_{\\rm X} \\gtrsim 10^{37}$ \\ergss) and high spin-up rates  (i.e., a significant increase in pulse frequency) and (b) normal outbursts (or type I), characterized by lower luminosities ($L_{\\rm X} \\sim 10^{36}-10^{37}$ \\ergss), low spin-up rates (if any), and recurrence at the orbital period \\citep{Stella86,Bildsten97}. As a population Be/X-ray binaries show a correlation between their spin and orbital periods \\citep{Corbet86,Waters89}.  EXO 2030+375 is a 42-second transient accreting X-ray pulsar discovered during a giant outburst in 1985 with {\\em EXOSAT} \\citep{Parmar89}. Optical and infrared observations of the {\\em EXOSAT} error circle identified a B0 Ve star as the most likely companion \\citep{Motch87,Janot88,Coe88}. The initial outburst was first detected at a 1-20 keV luminosity of $1 \\times 10^{38}$ \\ergss\\ (using a distance of 5 kpc assumed by Parmar et al. 1989) on 1985 May 18 and declined to $\\lesssim 3.8 \\times 10^{34}$ \\ergss\\ by 1985 August 25.  During this luminosity decline, the intrinsic spin period changed dramatically, with a characteristic spin-up timescale $-P/\\dot P \\approx 30$ yr \\citep{Parmar89}. This large intrinsic spin-up suggested that an accretion disk was present and made determination of an orbit difficult, resulting in 3 acceptable orbits.  The rate of change of pulse period $\\dot P$ \\citep{Parmar89}, the energy spectrum \\citep{Reynolds93,Sun94} and the 1-10 keV pulse profile \\citep{PWS89} all showed significant luminosity dependence. Further evidence of an accretion disk resulted from the detection of 0.2 Hz quasi-periodic oscillations \\citep{Angelini89} consistent with the magnetospheric beat frequency model \\citep{Lamb85} and the Keplerian frequency model \\citep{vanderKlis87}. A second outburst, roughly a factor of 10 weaker in luminosity than the first, was also detected with {\\em EXOSAT} in 1985 October 28-November 3 \\citep{Parmar89}. This outburst was apparently a normal outburst, but it showed unusual flaring activity. These flares were detected on 1985 October 30-31 and had a 4 hour recurrence period \\citep{Parmar89,Apparao91}. Observations of EXO 2030+375 were quite sparse for the next several years. {\\em Ginga} observed EXO 2030+375 on 1989 October 29-31 and 1991 October 24 \\citep{Sun92}. Pulsations were detected only in the 1989 observations, with an observed pulse period of 41.68202(8) s and an observed period derivative of $\\dot P = -(8.3 \\pm 0.9) \\times 10^{-9}$ s/s at MJD 47828.95. EXO 2030+375 was also detected in the soft X-ray band for two days near a periastron passage with {\\em ROSAT} in November 1990 \\citep{Mavro94}; however the observations, short 10-28 s scans, were not suitable to detect pulsations.  The most extensive observations of EXO 2030+375 were made with the Large Area Detectors (LADs) of the Burst and Transient Source Experiment \\citep[BATSE]{Fishman89} on the {\\em Compton Gamma Ray Observatory (CGRO)}. From launch in April 1991 until {\\em CGRO} was de-orbited in June 2000, BATSE provided nearly continuous coverage of EXO 2030+375. During the interval 1992 February 8-1993 August 26, 13 consecutive outbursts of EXO 2030+375 were seen with peak luminosities of $0.3 \\times 10^{37} \\leq L_{\\rm X\\ 1-20\\ keV} \\leq 3.0 \\times 10^{37}$ \\ergss\\ (See Stollberg et al.\\ 1999 for spectral assumptions), durations of 7-19 days, and spacings of 46 days \\citep{Bildsten97,Stollberg99}. These outbursts peaked 5-6 days after periastron passage. A few detections of marginal statistical significance preceeded and followed this sequence of outbursts. During the 13 outbursts, the pulsar spun-up at an average rate of $1.3 \\times 10^{-13}$ \\hzs. The 20-160 keV pulse profiles were double peaked and showed no significant energy or luminosity dependence. A binary orbit listed in the top row of Table~\\ref{tab:orb} was determined using these 13 outbursts \\citep{Stollberg99}. Near simultaneous optical, infrared, and X-ray observations of an outburst in 1993 June/July showed no significant evidence for a correlation between X-ray flux and infrared luminosity or between the X-ray flux and the equivalent width, strength, or profile of the H$\\alpha$ emission line \\citep{Norton94}. EXO 2030+375 was not detected above a flux level of $\\sim 4.2 \\times 10^{-10}$ \\ergcms\\ (20-50 keV) in the search techniques used from 1993 August 26 to 1996 April. Beginning in 1996 April, three short, weak, outbursts were detected with BATSE, separated by 46 days \\citep{Stollberg96,Stollberg99}. These outbursts began about 5 days prior to periastron passage as predicted using the orbital model of \\citet{Stollberg99}. Outbursts were also regularly detected with the All-Sky Monitor \\citep[ASM]{Levine96} on the {\\em Rossi X-ray Timing Explorer (RXTE)} beginning in 1996 March \\citep{ReigCoe98}.  An outburst of EXO 2030+375 was observed from 1996 July 1-10 with the {\\em RXTE} Proportional Counter Array \\citep[PCA]{Jahoda96}. Pulsations were detected throughout the observations. This outburst began about 5 days prior to periastron passage. The 2-10 keV pulse profile did not show significant intensity dependence. In fact, it was consistent with the 1-10 keV {\\em EXOSAT} profile observed at a similar luminosity \\citep{ReigCoe98}. The energy spectrum was correlated with luminosity \\citep{Reig99}. This correlation was consistent with an extrapolation of that observed in {\\em EXOSAT} data by \\citet{Reynolds93}. Evidence for a possible cyclotron feature at 36 keV was found in spectra from the High Energy X-ray Timing Experiment \\citep[HEXTE]{Rothschild98} on {\\em RXTE} \\citep{Reig99}.  In this paper we will present an improved orbit determination for EXO 2030+375 using BATSE, {\\em RXTE}, and {\\em EXOSAT} data. This improved orbit along with more sensitive search techniques has led to the detection of pulsations in 52 outbursts in 9 years of BATSE data, including several outbursts in the period 1993 August to 1996 April when the source was previously believed to be quiescent. Evidence for 19 additional outbursts, including 10 missed with BATSE and 9 after {\\em CGRO} was de-orbited, was observed with the {\\em RXTE} ASM, for a total of 71 outbursts observed out of 82 periastron passages from 1991 April to 2001 August. We show that our improved orbital parameters remove the two intervals of enhanced spin-up observed by \\citet{Reynolds96} in the initial {\\em EXOSAT} outburst, when they used the orbit of \\citet{Stollberg94} to correct the observed pulse periods. We also show evidence for a correlation between spin-up rate and flux in the BATSE data and compare it to that observed in the {\\em EXOSAT} data. We compare our X-ray results to optical and infrared observations and discuss evidence for a decline in the density of the Be disk and its effects on the X-ray flux and pulsed frequency histories. We discuss evidence in both optical observations of H$\\alpha$ profiles and X-ray observations that suggests a global one-armed oscillation (i.e., a density perturbation) was propagating in the Be disk. Lastly, we discuss our observations of EXO 2030+375 in context of current models. ",
        "conclusions": "\\subsection{Improvements to Sensitivity and Orbital Parameters}  Improvements to our techniques used to search for pulsations in the BATSE data reduced systematic errors to below statistical levels. These improvements included using a modified $Z_n^2$ statistic which accounted for aperiodic noise from either the measured source or others in the field of view such as Cygnus X-1 and automatically fitting Earth occultation steps from bright sources using a database including source locations, dates of activity, and flux levels to determine which sources needed to be fitted. These improvements allowed us to use much longer integrations for searches, e.g. 4 days, with accurately determined errors. Searches of BATSE data using these new techniques resulted in detection of 52 outbursts including several in the 2.5 year period from 1993 August to 1996 April when EXO 2030+375 was previously believed to be quiescent. Our results show that EXO 2030+375 has undergone an outburst near most likely every periastron passage for 9 years. From MJD 50643-51004 (1997 Jul - 1998 Jul), EXO 2030+375 went undetected with BATSE for the longest interval, 7 periastron passages. However, {\\em RXTE} PCA observations detected an outburst from EXO 2030+375 on the third periastron passage missed with BATSE and the {\\em RXTE} ASM detected outbursts for 6 of those periastron passages, suggesting that outbursts were still occurring but were below BATSE's detection threshold. Cygnus X-1 was noisy during this time interval. Additional noise plus a slight decrease in the intensity of EXO 2030+375 likely explains why BATSE missed these outbursts.  The 13 consecutive outbursts used by \\citet{Stollberg99} to determine an orbit for EXO 2030+375 provided good coverage from periastron passage to about 14 days after periastron passage. Their orbit fitting was further complicated by the intrinsic spin frequency variations during these 13 brighter outbursts. Comparing panel (b) in Figure~\\ref{fig:freqs} and Figure~\\ref{fig:goao} shows where there were problems with the \\citet{Stollberg99} orbital parameters. The outbursts from MJD 50000-50700 (1995 Oct - 1997 Sep) covered earlier orbital phases, from about 7.5 days before until 6 days after periastron passage, than those included in the orbit fitting of \\citet{Stollberg99}. These outbursts showed excess scatter in the spin frequencies determined using the orbital parameters of \\citet{Stollberg99}, indicating that fitting those outbursts would improve the orbital parameters. Fits to BATSE and {\\em RXTE} data from these outbursts in addition to the outbursts fitted by \\citet{Stollberg99} resulted in improved orbital parameters that removed the excess scatter in the spin frequencies. The orbital parameters were further improved by including pulse frequencies determined from {\\em EXOSAT} ME data during the initial giant outburst, which provided coverage of the entire orbit including two periastron passages with good coverage of orbital phases from 21 to 7 days before periastron passage. Fits to a broad range of orbital phases allowed us to decouple orbital and intrinsic effects and to provide a good determination of the orbital parameters without using complicated models for the intrinsic torque variations. Our new orbital parameters are given in Table~\\ref{tab:orb}.  \\subsection{Orbital Phasing of Outbursts}  The outbursts prior to MJD 50000 (1995 Oct) peaked at a very regular orbital phase of about 6 days after periastron passage (see Figure~\\ref{fig:goao}). The outburst just after MJD 50000 peaked at a much earlier orbital phase, 4 days before periastron passage and then gradually recovered to peak at a new stable orbital phase of about 2.5 days after periastron passage. A possible explanation of the shift in orbital phase of the outbursts is a density perturbation (global one armed oscillation) in the Be disk. Evidence for these density perturbations is seen in the H$\\alpha$ line profiles for several Be/X-ray binaries \\citep[for example]{Negueruela01a, Negueruela01b}. When a density perturbation is present, the H$\\alpha$ line is double peaked. The relative size of the two peaks changes with a cycle of several years. The density perturbation produces a non-axially symmetric Be disk. If the pulsar interacts with a region of the disk affected by the perturbation, more material would be available for accretion, possibly causing a shift in outburst phase. The trend in the orbital phases from MJD 50000-50600 (1995 Oct - 1997 Jun) in Figure~\\ref{fig:goao} has a slope of $m = 0.0085 \\pm 0.0017$. This slope can be expressed in terms of a beat frequency between the orbital period and some other period. Assuming the other period is longer than the orbital period, as is expected for global one-armed oscillations, then \\begin{equation} m = \\frac{\\nu_{\\rm beat}}{\\nu_{\\rm orb}}-1 \\end{equation} where $\\nu_{\\rm beat}$ is the beat frequency and $\\nu_{\\rm orb}$ is the orbital frequency. If this trend is due to beating between the orbital period and the period of the density perturbation, then the beat frequency of is given by \\begin{equation} \\nu_{\\rm beat} = \\nu_{\\rm orb} + \\nu_{\\rm perturb} \\end{equation} where $\\nu_{\\rm perturb}$ is the frequency of the density perturbation. Solving for $\\nu_{\\rm perturb}$, we get a perturbation frequency of $(1.8 \\pm 0.4) \\times 10^{-4}$ cycles day$^{-1}$ or equivalently, a period of $15 \\pm 3$ years for the density perturbation to propagate around the Be disk. No significant change in X-ray intensity was seen when the outbursts shifted in orbital phase, although 3 outbursts went undetected near the time when the shift occurred. {\\em EXOSAT} observations of the second outburst of EXO 2030+375 after its discovery (1985 October 29-November 3, MJD 46367-46372) show that the outburst peaked at $9.5 \\pm 1.1$ days after periastron, assuming that the peak flux detected with {\\em EXOSAT} is the peak of a normal outburst. Interestingly, this is 3.5 days later than the peak time observed in the pre-MJD 50000 (1995 Oct) BATSE data, which is also 3.5 days later than the peak time observed after MJD 51000 (1998 Jul), when the peak time had stopped changing rapidly and occurred about 2.5 days after periastron. This suggests that another shift in outburst phase may have occurred between the {\\em EXOSAT} and BATSE observations. However, if such a shift occurred, it suggests a shorter propagation period of $\\lesssim 10$ years.  Optical observations of the H$\\alpha$ profile (Figure~\\ref{fig:ha}) clearly indicate that the structure of the circumstellar disk around the Be star changed significantly at some time between the 1996 July (MJD 50273) observation and the 1997 August (MJD 50661) observation. These observations support the idea of a global one-armed oscillation propagating in the Be disk suggested by the shift in orbital phase of the outbursts. However, sparse observations of the H$\\alpha$ profile do not allow us to directly correlate H$\\alpha$ profile changes with changes in the X-ray outbursts.  \\subsection{Relationship between X-ray and IR Measurements\\label{sec:xir}}  Comparing the IR measurements in Figure~\\ref{fig:hjk} to the X-ray measurements (Figures~\\ref{fig:freqs} and \\ref{fig:longterm}) suggests a relationship between the IR behavior and the X-ray activity. The IR measurements indicate that the Be disk was fairly stable and roughly constant in density from near the end of the initial giant outburst until MJD 49000 (1993 Jan). Figure~\\ref{fig:longterm} indicates that the pulsar was spinning up for most of this period. Near MJD 49000, the density of the disk began to decline. Since the disk was becoming less dense, the reservoir of material available to the pulsar near periastron passage was slowly reduced. After MJD 49250 (1993 Sep), the  X-ray pulsed flux responded to the lower density disk and dropped dramatically. Peak 20-50 keV pulsed fluxes dropped from $(3-5) \\times 10^{-10}$ \\ergcms\\ to  $\\lesssim 1.5 \\times 10^{-10}$ \\ergcms. The global spin-up rate of the pulsar took longer to respond. After about MJD 49250, the spin-up rate slowed and by MJD  49400 (1994 Feb), the pulsar had begun to spin-down. No obvious response in the X-rays was seen to the slow increase in density of the Be disk indicated by the IR measurements after MJD 50300 (1996 Aug). Perhaps the disk had not yet become dense enough to increase the mass accretion rate to a level where the pulsar would begin to spin-up.  To determine whether or not the observed spin-down was likely due to centrifugal inhibition of accretion \\citep{Stella86}, i.e., the propeller effect \\citep{Illarionov75}, we estimate the flux at the onset of this effect by equating the magnetospheric radius to the corotation radius. The magnetospheric radius is given by \\citep{Pringle72, Lamb73} \\begin{equation} r_{\\rm m} \\simeq k (G M)^{1/7} \\mu^{-2/7} L^{-2/7} R^{-2/7}\\label{eqn:rm} \\end{equation} where $G$ is the gravitational constant; $M$ and $R$ are the mass and radius of the neutron star; and $L$ is the luminosity. $k$ is a constant factor of order 1. Equation~\\ref{eqn:rm} with $k \\simeq 0.91$ gives the Alfv\\'en radius for spherical accretion and with $k \\simeq 0.47$ gives the magnetospheric radius derived by \\citet{Ghosh79}. The corotation radius is given by \\begin{equation} r_{\\rm co} = (G M)^{1/3} (2 \\pi \\nu)^{-2/3}\\label{eqn:rco} \\end{equation} where $\\nu$ is the spin frequency of the pulsar. Setting $r_{\\rm m} = r_{\\rm co}$ gives the threshold flux for the onset of centrifugal inhibition of accretion, i.e., \\begin{equation} F_{\\rm x}^{\\rm min} \\simeq 3 \\times 10^{-10}\\ {\\rm ergs}\\ {\\rm cm}^{-2}\\ {\\rm s}^{-1}\\ k^{7/2} \\mu_{30}^2 M_{1.4}^{-2/3} R_6^{-1} P_{\\rm 41.7s}^{-7/3} d_{\\rm kpc}^{-2}\\label{eqn:cia} \\end{equation} where $\\mu_{30}$, $M_{1.4}$, $R_6$, and $P_{\\rm 41.7 s}$ are the pulsar's magnetic moment in units of $10^{30}$ G cm$^{3}$, mass in units of 1.4 $M_{\\odot}$, radius in units of $10^6$ cm, and spin period in units of 41.7 seconds, respectively. \\citet{Reynolds96} fitted 3 accretion torque models, a simple spherical accretion model, the \\citet{Ghosh79} model, and the \\citet{Wang87} model to the giant outburst to estimate values for $\\mu_{30}$ and $d_{\\rm kpc}$. For the simple spherical accretion model, where $k \\simeq 0.91$, \\citet{Reynolds96} obtained $\\mu_{30} \\simeq 5$ for an assumed distance of $d_{\\rm kpc} = 5$. Both the \\citet{Ghosh79} model and the \\citet{Wang87} model use $k \\simeq 0.47$. \\citet{Reynolds96} obtained $\\mu_{30} \\simeq 12$ and $d_{\\rm kpc} \\simeq 5.2$ from fits to the \\citep{Ghosh79} model. \\citet{Parmar89} also fitted the \\citet{Ghosh79} model, obtaining $\\mu_{30} \\simeq 20$ and $d_{\\rm kpc} \\simeq 5.3$ for the first nine period measurements and $\\mu_{30} \\simeq 11$ and $d_{\\rm kpc} \\simeq 5.0$ for the first 10 period measurements. \\citet{Reynolds96} fit of the \\citet{Wang87} model yielded the lowest distance $ d_{\\rm kpc} \\simeq 4.1$ with $\\mu_{30} \\simeq 10$.  Substituting \\citet{Reynolds96} values yields $F_{\\rm x}^{\\rm min} \\simeq (1.1-2.2) \\times 10^{-10}$ \\ergcms\\ while \\citet{Parmar89}'s values yield $F_{\\rm x}^{\\rm min} \\simeq (1-3) \\times 10^{-10}$ \\ergcms.  The minimum flux where pulsations were first detected during the faint 1996 July outburst observed with {\\em RXTE} was $3.3 \\times 10^{-10}$ \\ergcms\\  \\citep[2.7-30 keV]{Reig99}, which is comparable to the lowest flux of $4 \\times 10^{-10}$ \\ergcms\\ observed by \\citet{Parmar89} before pulsations became undetectable in the giant outburst. Since pulsations were detected with {\\em RXTE} and no significant changes in the 2-10 keV pulse profile were observed relative to higher fluxes \\citep{ReigCoe98}, this flux is most likely above the threshold for centrifugal inhibition of accretion, which is consistent with our calculations in the previous paragraph. On 1985 August 25, when pulsations were not detected, \\citet{Parmar89} measured an upper limit of $1.3 \\times 10^{-11}$ \\ergcms\\ on the flux from EXO 2030+375, which places a lower bound on $F_{\\rm x}^{\\rm min}$, that is also consistent with our calculations, assuming that the observed sudden drop in flux was due to centrifugal inhibition of accretion.  \\subsection{Spin-up vs. flux correlations}  In giant outbursts of Be/X-ray binaries, accretion disks are expected to be present and indeed, evidence for an accretion disk, based on a correlation between the observed flux and spin-up rate, has been found for several sources including EXO 2030+375 \\citep{Parmar89, Reynolds96, Stollberg99} during giant outbursts of Be/X-ray binaries \\citep{Wilson98, Bildsten97}. Independent evidence for an accretion disk based on the detection of quasi-periodic oscillations during a giant outburst has been found for EXO 2030+375 \\citep{Angelini89} and A0535+262 \\citep{Finger96b}. In addition to the expected correlation between observed flux and spin-up rate in the giant outburst of EXO 2030+375, the BATSE data suggest a correlation is also present in the brighter normal outbursts of this system (See Figure~\\ref{fig:fdotvsflux}). Until recently, normal outbursts were believed to be due to direct wind accretion from the Be disk, so significant spin-up was not expected because wind accretion is not believed to be very efficient at transferring angular momentum \\citep{Ruffert97}. However, evidence for spin-up during normal outbursts has been observed in GS 0834--430 \\citep{Wilson97}, 2S 1417--624 \\citep{Finger96a}, 2S 1845--024 \\citep{Finger99}, and previously in EXO 2030+375 \\citep{Stollberg99}. In BATSE observations of 2S 1845--024, a correlation between the spin-up rate and the pulsed flux was also observed.  The comparison of the BATSE and {\\rm EXOSAT} data in Figure~\\ref{fig:fdotvsflux}, despite an uncertain bolometric correction, shows that at the lower fluxes measured with BATSE the correlation falls off more rapidly than a power law. This suggests either that spin-down torques become important or perhaps that a disk forms during the outburst and we are averaging over periods of wind and disk accretion. The data are clearly inconsistent with the power law index of $6/7$ predicted from simple accretion theory, which does not consider spin-down torques. The \\citet{Ghosh79} model fitted by \\citet{Reynolds96}, which assumes an accretion disk is present, roughly follows the trend observed in the data; however, the brightest {\\em EXOSAT} observations, which drive the fit, clearly deviate from the model. The observed flux is given by $F = (\\beta L)/(4 \\pi d^{2})$, where $\\beta$ is a beaming factor.  This beaming factor, which is typically assumed to be equal to one for simplicity, depends on the pattern of emitted radiation at the pulsar and cannot be determined without modeling of the pulse profiles (which is beyond the scope of this work). Because large luminosity dependent 1-10 keV pulse profile variations were observed during the giant outburst \\citep{PWS89}, one would expect that the beaming factor was also changing with luminosity. To fit any of the discussed models, the beaming factor must decrease with increasing luminosity.  The spin-up rate and its correlation with pulsed flux during the earlier outbursts of EXO 2030+375 observed with BATSE suggest an accretion disk may be present. A disk will form if the specific angular momentum of the material accreted from the Be star's disk is comparable to the Keplerian specific angular momentum at the magnetospheric radius. The specific angular momentum $\\ell$ of the accreted material is given by \\begin{equation} \\ell = 2 \\pi I \\dot \\nu \\dot M^{-1} \\simeq (4.8-9.1) \\times 10^{16}\\ \\rm{cm}\\ \\rm{s}^{-1}\\ d_{\\rm 7.1 kpc}^{-2},\\label{eqn:ell} \\end{equation} assuming $\\dot M = L (G M/ R)^{-1}$. Here $\\dot \\nu = 6.5 \\times 10^{-13}$ Hz s$^{-1}$ is the spin-up rate, $\\dot M$ is the mass accretion rate, $L = (0.8-1.6) \\times 10^{37}$ \\ergss\\ $d_{\\rm 7.1 kpc}^2$ is the luminosity, and $d_{\\rm 7.1 kpc}$ is the distance in units of 7.1 kiloparsecs. The spin-up rate and luminosity are outburst averaged values from a typical bright outburst observed with BATSE in 1992 February. The luminosity is estimated from BATSE pulsed flux of $2.5 \\times 10^{-10}$ \\ergcms\\ using a bolometric correction of $8.0 \\pm 2.4$ (See Section~\\ref{sec:torq}), and a distance of 7.1 kpc (See Section~\\ref{sec:oir}). We have assumed typical pulsar parameters, listed in Section~\\ref{sec:xir}. If $\\ell \\sim \\ell_{\\rm m} = (GMr_{\\rm m})^{1/2}$, the Keplerian specific angular momentum at the magnetospheric radius, a disk will form. Once a disk has formed the specific angular momentum of accreting material is maintained near $\\ell_{\\rm m}$. At $F_{\\rm x}^{\\rm min}$ (Equation~\\ref{eqn:cia}), $r_{\\rm m} = r_{\\rm co}$. Hence $\\ell_{\\rm m}$ can be expressed as a function of $F_{\\rm x}^{\\rm min}$ and $r_{\\rm co}$, which is independent of the distance and magnetic field of the pulsar. \\begin{equation} \\ell_{\\rm m} \\simeq 6.1 \\times 10^{17}\\ {\\rm cm}^2\\ {\\rm s}^{-1}\\ P_{\\rm 41.7s}^{1/3} \\left(\\frac{F}{F_{\\rm x}^{\\rm min}}\\right)^{-1/7} \\end{equation} In the previous section, using {\\em RXTE} and {\\em EXOSAT} observations, we determined $1.3 \\times 10^{-11}$ \\ergcms\\ $\\lesssim F_{\\rm x}^{\\rm min} \\lesssim 3.3 \\times 10^{-10}$ \\ergcms.  With the estimated average bolometric flux of $(1.4-2.6) \\times 10^{-9}$ \\ergcms\\ from the 1992 February outburst used in Equation~\\ref{eqn:ell}, this yields $\\ell_{\\rm m} \\simeq (2.9-5.0) \\times 10^{17}$ cm$^2$ s$^{-1}$, which is within an order of magnitude of $\\ell \\simeq (4.8-9.1) \\times 10^{16}$ cm$^2$ s$^{-1}$, suggesting a disk is likely present because considerable angular momentum is present in the system. If a disk forms during the outburst, our average values of $\\dot \\nu$ and $F$ used in these calculations would include periods of wind accretion and periods of disk accretion, possibly explaining why $\\ell \\simeq (0.1-0.3) \\ell_{\\rm m}$ rather than being a larger fraction. In contrast, for the wind-fed system Vela X-1 where a disk is not expected to be present, $\\dot \\nu \\simeq 6 \\times 10^{-14}$ Hz s$^{-1}$ \\citep{Inam00}, $L \\simeq 2 \\times 10^{38}$ \\ergss, and $\\mu \\simeq 2.1 \\times 10^{30}$ G cm$^{3}$ \\citep{Makishima99} leading to $\\ell \\simeq 3.5 \\times 10^{14}$ cm$^2$ s$^{-1}$ and $\\ell_{\\rm m} \\simeq 2.2 \\times 10^{17}$ cm$^2$ s$^{-1}$, i.e., $\\ell \\simeq 0.002 \\ell_{\\rm m}$.  \\subsection{Current Models}  Although the direct wind accretion model has difficulty explaining evidence for an accretion disk during normal outbursts, the viscous decretion disk model \\citep{Lee91, Porter99, Okazaki01a} provides a natural explanation. In this model, the inner edge of the Be disk has a Keplerian velocity. Viscosity conducts material outwards, so that it moves in quasi-Keplerian orbits with low radial velocities. The radial outflow is subsonic for the orbital sizes of all Be/X-ray binaries with a known solution. This model successfully accounts for most observations of Be disks \\citep{Okazaki01}. \\citet{Negueruela01a} found that tidal interaction of the neutron star truncates the circumstellar Be disk. In the disk, truncation occurs when the outward viscous torque is less than the inward resonant torque which truncates the disk at a resonant radius. Because of the truncation, the Be disk cannot reach a steady state \\citep{Okazaki01}. According to the modeling of \\citet{Okazaki01}, the Be disk in EXO 2030+375 is likely truncated at the 4:1 resonance radius, which is close to the radius of the critical lobe at periastron. If the truncation radius is close to or slightly beyond the critical lobe radius at periastron, material with high angular momentum will flow through the first Lagrangian point to the neutron star, making formation of a transient accretion disk likely \\citep{Okazaki01} in a normal outburst. Hence this model provides a reasonable explanation of the brighter outbursts.  We propose the following scenario to explain the IR and X-ray observations. The Be disk is truncated at a 4:1 resonance radius. Following \\citet{Okazaki01}, there should be an outburst at every periastron unless the circumstellar disk disappears. Around MJD 49000 (sometime in 1993), a major structural change occurred in the Be star's circumstellar disk. It became much less dense, as shown by the change in IR magnitudes and the H$\\alpha$ equivalent width. Much less matter was available for accretion, and as a consequence, the X-ray flux dropped and the neutron star spin-up ended. At the same time, or shortly afterward, a density wave (or global one-armed oscillation) developed and began to precess, without interacting with the neutron star's orbit. Around MJD 50000 (late in 1995), the density perturbation interacted with the neutron star's orbit, at a phase corresponding to about 3 days prior to periastron passage, producing an X-ray outburst peaked at that phase. This implies that the precession of the density perturbation was prograde, in the same sense as the neutron star's orbital motion. Sometime after about MJD 50600 (mid-1997), the density perturbation lost contact with the neutron star's orbit, in a position symmetrical with respect to periastron, i.e., about 3 days after periastron. This ended the fast migration of the outburst peaks in orbital phase. Some slow migration may still be present in Figure~\\ref{fig:goao}, which may eventually lead to the previous value of 6 days after periastron, after a complete precession period of the perturbed disk."
    },
    "0201/astro-ph0201011_arXiv.txt": {
        "abstract": "{The relation between the Ultra Compact Objects (hereafter UCOs) recently discovered in the Fornax cluster (Drinkwater et al. \\cite{Drinkw00a}, Hilker et al. \\cite{Hilker99}) and the brightest globular clusters associated with the central galaxy NGC 1399 has been investigated. The question was adressed whether the UCOs constitute a distinct population of objects not linked to globular clusters or whether there is a smooth transition between both populations.\\\\ Therefore, a spectroscopic survey on compact objects in the central region of the Fornax cluster was carried out  with the 2.5m du Pont telescope (LCO). UCOs and the bright NGC 1399 globular clusters with similar brightness were inspected. 12 GCs from the bright end of the globular cluster luminosity function have been identified as Fornax members. Eight are new members, four were known as members from before. Their magnitude distribution supports a smooth transition between the faint UCOs and the bright globular clusters. There is no evidence for a  magnitude gap between both populations. However, the brightest UCO clearly stands out, it is too bright and too large to be accounted for by globular clusters. For the only UCO included in our survey, a relatively high metallicity of $[\\frac{Fe}{H}]\\simeq -0.5$ dex is measured.\\\\ ",
        "introduction": "\\subsection{Magnitude - surface brightness relation of early type dwarf galaxies}  The faint end of the galaxy luminosity function is mainly populated by dwarf elliptical galaxies (dEs) and dwarf spheroidals (dSphs, the faintest dEs in the Local Group). These galaxies are the most numerous type of galaxies in the nearby universe, having absolute magnitudes fainter than $M_{V}$ $\\simeq$ $-$17 mag. They follow a tight magnitude-surface brightness relation in the sense that central surface brightness increases with increasing luminosity (Ferguson \\& Sandage \\cite{FergSan88}, \\cite{FergSan89}). The validity of this relation has been a subject of lively debate over the last decade. A number of  authors have argued against the existence of a magnitude-surface brightness relation for dEs (Davies et al. \\cite{Davies88}, Phillipps et al. \\cite{Philli88}, Irwin et al. \\cite{Irwin90}) and questioned the cluster membership assignement to dEs based on morphology.\\\\ Recently, spectroscopic membership confirmation has shed light into this matter. Drinkwater et al. (\\cite{Drinkw01a}) do confirm the brightness - surface brightness relation for Fornax dwarfs based on the data of their Fornax Cluster Spectroscopic Survey (FCSS). They obtain spectra for all objects in the central 2 degrees of the Fornax cluster down to a limiting magnitude of $M_V\\sim-12.5$ mag and find that the magnitude surface brightness relation for dEs is well defined. Under this scheme, galaxies whose total luminosity and/or surface brightness lie significantly outside this relation are hard to classify. Two of the few examples for such peculiar objects that have been known for a long time are the high surface brightness compact dwarf elliptical (cdE) M32 and the very extended low surface brightness spiral Malin 1 (Bothun et al. \\cite{Bothun87}, Impey et al. \\cite{Impey88}, Bothun et al. \\cite{Bothun91}).\\\\  \\subsection{New Ultra Compact Objects} Most recently, in the course of the FCSS, Drinkwater et al. (\\cite{Drinkw00a}) detected five ultra compact objects (UCOs, UCDs in their papers) within 30$'$ projected distance from the Fornax cluster's central galaxy, NGC 1399\\footnote{Two of them had already been detected by Hilker et al. (\\cite{Hilker99})}. Although as bright as average size dEs, they are by far more compact. Four of the five UCOs have absolute magnitudes of about $M_V = -12$ mag, one is significantly brighter with $M_V = -13.3$ mag. On HST-STIS images, the four fainter UCOs have King profile effective radii between 10 and 17 pc, while the brightest one has about 50 pc (Drinkwater et al. \\cite{Drinkw01b}). In Table \\ref{ucocor}, the properties of the UCOs are summarized. All of them are significantly brighter than the brightest galactic globular cluster ($\\omega$ Centauri has $M_V=-10.2$ mag) but significantly fainter than M~32 ($M_V=-16$ mag). Neither Hilker et al. (\\cite{Hilker99}), Drinkwater et al. (\\cite{Drinkw00a}, \\cite{Drinkw01b}) nor Phillipps et al. (\\cite{Philli01}) could draw definite conclusions about the nature of the UCOs.\\\\ \\begin{table*} \\caption{\\label{ucocor} Properties of the five UCOs. Adopted distance modulus to Fornax 31.3 mag. Magnitudes are from a photometric wide field survey of the central Fornax cluster (Hilker et al. \\cite{Hilker02}, in prep.). $^*$ UCO~3 and 4 were detected first by Hilker et al. (\\cite{Hilker99}).} \\begin{flushleft} \\begin{center} \\begin{tabular}{llllllll} \\hline\\noalign{\\smallskip} Name & $\\alpha$ (2000) & $\\delta$ (2000) &$V$ [mag] & $M_V$ [mag]&$(V-I)$ [mag]&$r_{\\rm eff}$ [pc]&$d$(NGC 1399) [$'$]\\\\\\hline\\hline UCO~1 & 03:37:03.30 & -35:38:04.6 & 19.31 & -11.99 & 1.17&12&20.74\\\\ UCO~2 & 03:38:06.33 & -35:28:58.8 &19.23 & -12.07 & 1.14&15&5.10\\\\ UCO~3$^*$ & 03:38:54.10 & -35:33:33.6 &18.06 & -13.24 & 1.20&50&8.26\\\\ UCO~4$^*$ & 03:39:35.95 & -35:28:24.5 &19.12 & -12.18 & 1.14&17&13.64\\\\ UCO~5 & 03:39:52.58 & -35:04:24.1 &19.50 & -11.80 & 1.04&10&28.26\\\\ \\hline \\end{tabular} \\end{center} \\end{flushleft} \\end{table*}One possibility is that UCOs are bright globular clusters. NGC 1399 has a very rich globular cluster system with about 6000 GCs within 10$'$ (about 40 kpc at Fornax distance) from its center (Kohle et al. \\cite{Kohle96}, Forbes et al. \\cite{Forbes98}) which may contain GCs as bright as the UCOs. Another possibility is that UCOs are the nuclei of stripped dwarf galaxies. Threshing dE,Ns in the cluster potential has been shown to work (Bekki et al. \\cite{Bekki01}). Lotz et al. (\\cite{Lotz01}) find that the luminosity function of 27 Fornax and Virgo dwarf nuclei peaks at $V=21.7$ mag ($M_{V}=-9.6$ mag) with a dispersion of $\\sigma=1.2$ mag, whereas the Globular Cluster Luminosity Function (GCLF) peaks at about $M_V=-7.4$ mag. So, nuclei are on average more than 2 mag brighter than GCs and might mix up with the bright tail of the GCLF. Or do the UCOs represent a new group of ultra compact dwarf galaxies, extreme cases of M~32? No discriminating statement could be made until now.\\\\ \\noindent It would be very interesting to know whether the UCOs had an origin different from the population of globular clusters of NGC 1399.\\\\ The UCOs have magnitudes roughly equal to the completeness magnitude limit of the FCSS ($V=19$ mag or $M_V=-11.7$ mag). This is about 4.5 magnitudes brighter than the turnover of the globular cluster luminosity function (Kohle et al. \\cite{Kohle96}). In Fig.~\\ref{introlf}, the three relevant luminosity functions are shown in one plot: the LF of all observed sources in Drinkwater et al.'s FCSS (to indicate its completeness limit); the LF of the UCOs; and the bright end of NGC 1399's GCLF, respresented by a Gaussian with $V_{\\rm to}=23.9$ mag and $\\sigma=1.2$ mag, as taken from Kohle et al.'s GCLF. A total number of 8100 GCs was adopted, which is the number contained within 20$'$ from NGC 1399 (see Sect.~\\ref{totalnum}). \\\\ \\begin{figure}[h!] \\vspace{-0.5cm} \\epsfig{figure=mieskef1.eps,width=8.6cm,height=8.6cm} \\caption{\\label{introlf}{\\it Solid line}: GCLF of NGC 1399, represented by a Gaussian with $V_{\\rm to}=23.9$ mag and $\\sigma=1.2$ mag. Both values taken from Kohle et al. (\\cite{Kohle96}). {\\it Solid histogram}: LF of all observed sources in the FCSS in the central Fornax cluster, divided by four to fit into the plot (Drinkwater et al. \\cite{Drinkw01c}, $V$ magnitudes from Hilker et al. \\cite{Hilker02}, in prep.). It is shown to illustrate the FCSS's completeness limit (90\\% at about $V=19$ mag). {\\it Short dashed histogram}: LF of the UCOs.} \\end{figure}\\\\ \\subsection{Aim of this paper} As one can see in Fig.~\\ref{introlf}, the magnitude range between the UCOs and the bright globular clusters must be probed more thoroughly, both  to know whether UCOs extend to fainter magnitudes and to determine the bright end of the GCLF. A gap in magnitude space between both populations would imply that the UCOs are very compact dEs or nuclei of stripped dwarfs rather than globular clusters.\\\\ In this paper, we describe and analyse a survey of compact objects in the central region of the Fornax cluster of galaxies that closes the magnitude gap between the FCSS and the globular cluster regime.\\\\ In Sect.~\\ref{Selcand} we describe the selection of candidates for the survey. In Sect.~\\ref{obsdatared} the observations and data reduction are described. Sect.~\\ref{measureradvel} shows the radial velocity measurement. In Sect.~\\ref{analysis} the results are analyzed, the question whether or not the UCOs are a distinct population is discussed and metallicities of a number of Fornax dE,Ns are measured. These results are discussed in Sect.~\\ref{discussion}. Finally, in Sect.~\\ref{summary} a summary and conclusions are presented.\\\\ ",
        "conclusions": "\\label{summary} In this paper, we presented a spectroscopic survey on compact objects in the central region of the Fornax cluster. The aim was to survey the magnitude regime framed by the newly discovered ultra compact objects (UCOs) in Fornax (Drinkwater et al. \\cite{Drinkw00a}, Hilker et al. \\cite{Hilker99}) and the brightest GCs around NGC 1399. In the FCSS, performed by Drinkwater et al., the UCOs are at the faint magnitude limit of their survey. We wanted to know whether these detections constitute the bright tail of a continous luminosity distribution with a smooth transition into the GC regime or whether they are a separate population.\\\\ The velocity measurements (cf. Sect.~\\ref{measureradvel}) resulted in the discovery of 12 GC candidates in the magnitude range between $19.70<V<20.95$ mag ($-11.6<M_{V}<-10.35$), located all within 20$'$ from NGC 1399. For four of them, cluster membership was known before. Their mean colour is $<V-I>$ = 1.11 $\\pm$ 0.11 mag and their mean radial velocity 1300 $\\pm$ 109 km/s with $\\sigma$=377 km/s.\\\\ In the subsequent analysis of the discoveries (cf. Sect.~\\ref{analysis}), the following results were obtained:\\\\ 1. The expected number of observed GCs originating from NGC 1399's GCS is 15 $\\pm$ 4, in good agreement with the 12 objects found. To calculate the expected number we assumed a Gaussian LF with $V_{\\rm to}$ = 23.9 mag and $\\sigma$ = 1.2 mag taken from Kohle et al. (\\cite{Kohle96}), used the GC surface density from Dirsch et al. (\\cite{Dirsch01}) and took into account photometric and geometric incompleteness.\\\\ 2. Assuming a more extended $t_5$ distribution for the LF as adopted by Kohle et al. (\\cite{Kohle96}), the expected number of observed GCs rises to 44 $\\pm$ 8, which is more than three $\\sigma$ higher than the number of observed objects. This implies that the LF has no extended bright wings. Hilker et al. (\\cite{Hilker99}) found that only by assuming such an extended LF the brightest UCO can be explained as a GC.\\\\ 3. There is no significant gap in magnitude space between our GC candidates and the four fainter UCOs within 20$'$ of NGC 1399. The GCS of NGC 1399 appears to extent to $M_{V}\\simeq -12$ mag.\\\\ 4. The only UCO included in our survey is slightly better fit by early stellar types, in contrast to the dE,Ns.\\\\ 5. The only UCO included in our survey has a relatively high metallicity compared to the dE,Ns, but is in the range of metal rich GCs or very compact dwarfs with almost solar metallicity like M~32.\\\\ Considering only point 5, we can neither rule out nor confirm that the UCOs are bright GCs. From points 1 to 4 we conclude: the four UCOs fainter than $V=19.1$ mag ($M_{V}=-12.2$ mag) can be well explained by the bright tail of the GCLF of NGC 1399. However, the apparent overlap of the two LFs is not sufficient to exclude the existence of stripped nuclei of formerly extended dE,Ns mixing up with the bright GCs. The faint UCOs are probably no extremely faint examples of cdEs, but are the brightest members of their object class (GCs or stripped nuclei). UCO~3, however, is so bright and large, that it probably is not a GC. It remains the most puzzling object.\\\\\\\\"
    },
    "0201/astro-ph0201294_arXiv.txt": {
        "abstract": "{ We analyse IRAS and COBE DIRBE data at wavelengths between 2.2 and 240\\,\\mim\\ of the central 500\\,pc of the Galaxy and derive the large-scale distribution of stars and interstellar matter in the Nuclear Bulge. Models of the Galactic Disk and Bulge are developed in order to correctly decompose the total surface brightness maps of the inner Galaxy and to apply proper extinction corrections. The Nuclear Bulge appears as a distinct, massive disk-like complex of stars and molecular clouds which is, on a large scale, symmetric with respect to the Galactic Centre. It is distinguished from the Galactic Bulge by its flat disk-like morphology, very high density of stars and molecular gas, and ongoing star formation. The Nuclear Bulge consists of an $R^{-2}$\\ Nuclear Stellar Cluster at the centre, a large Nuclear Stellar Disk with radius 230$\\pm$20\\,pc and scale height 45$\\pm$5\\,pc, and the Nuclear Molecular Disk of same size. \\newline The total stellar mass and luminosity of the Nuclear Bulge are 1.4$\\pm$0.6$\\times 10^9$\\,\\msun\\ and 2.5$\\pm$1$\\times 10^9$\\,\\lsun, respectively. About 70\\%\\ of the luminosity is due to optical and UV radiation from young massive Main-Sequence stars which are most abundant in the Nuclear Stellar Cluster. For the first time, we derive a photometric mass distribution for the central 500\\,pc of the Galaxy which is fully consistent with the kinematic mass distribution. We find that the often cited $R^{-2}$ distribution holds only for the central $\\sim$\\,30\\,pc and that at larger radii the mass distribution is dominated by the Nuclear Stellar Disk which has a flatter density profile. \\newline The total interstellar hydrogen mass in the Nuclear Bulge is $M_{\\rm H}$=2$\\pm$0.3$\\times 10^7$\\,\\msun, distributed in a warm inner disk with $R$=110$\\pm$20\\,pc and a massive, cold outer torus which contains more than 80\\%\\ of this mass. Interstellar matter in the Nuclear Bulge is very clumpy with $\\sim$90\\%\\ of the mass contained in dense and massive molecular clouds with a volume filling factor of only a few per cent. This extreme clumpiness, probably caused by the tidal stability limit in the gravitational potential of the Nuclear Bulge, enables the strong interstellar radiation field to penetrate the entire Nuclear Bulge and explains the relatively low average extinction towards the Galactic Centre. In addition, we find 3$\\times 10^7$\\,\\msun\\ of cold and dense material outside the Nuclear Bulge at positive longitudes and 1$\\times 10^7$\\,\\msun\\ at negative longitudes. This material is not heated by the stars in the Nuclear Bulge and gives rise to the observed asymmetry in the distribution of interstellar matter in the Central Molecular Zone. ",
        "introduction": "\\label{intro}  \\begin{table}[htb] \\caption[]{Abbreviations used in this paper  \\label{abbtab}}\\vspace{-0mm} \\begin{flushleft} \\begin{tabular}[t]{ll} \\hline \\noalign{\\smallskip} Abbreviation  & Meaning                       \\\\ \\noalign{\\smallskip} \\hline\\noalign{\\smallskip} ~~~~BH   & Black Hole                         \\\\ ~~~~BW   & Baade's Window                     \\\\ ~~~~CMZ  & Central Molecular Zone             \\\\ ~~~~FIR  & far-infrared ($\\sim${\\bf 30--300}\\,\\mim) \\\\ ~~~~FWHM & Full Width Half Maximum            \\\\ ~~~~GB   & Galactic Bulge                     \\\\ ~~~~GC   & Galactic Centre                    \\\\ ~~~~GD   & Galactic Disk                      \\\\ ~~~~GMC  & Giant Molecular Cloud              \\\\ ~~~~HPBW & Half Power Beam Width              \\\\ ~~~~HWHM & Half Width Half Maximum            \\\\ ~~~~ISM  & Interstellar Matter                \\\\ ~~~~KLF  & K-band Luminosity Function         \\\\ ~~~~MIR  & mid-infrared ($\\sim${\\bf 7--30})\\,\\mim)  \\\\ ~~~~MS   & Main Sequence                      \\\\ ~~~~{\\bf NB}  & {\\bf Nuclear Bulge}           \\\\ ~~~~{\\bf NMD} & {\\bf Nuclear Molecular Disk}  \\\\ ~~~~{\\bf NSC} & {\\bf Nuclear Stellar Cluster} \\\\ ~~~~{\\bf NSD} & {\\bf Nuclear Stellar Disk}    \\\\ ~~~~NIR  & near-infrared ($\\sim${\\bf 1--7}\\,\\mim)   \\\\ ~~~~PAH  & Polycyclic Aromatic Hydrocarbon    \\\\ ~~~~SED  & Spectral Energy Distribution       \\\\ ~~~~ZL   & Zodiacal Light                     \\\\ \\noalign{\\smallskip} \\hline\\noalign{\\smallskip} \\end{tabular} \\end{flushleft} \\end{table}  Our present knowledge of the Galactic Centre (GC) Region has recently been reviewed by, e.g., Blitz et al. (1993), Genzel et al. (1994), Mezger et al. (1996; hereafter MDZ96), and Morris \\& Serabyn (1996). The gas dynamics of the inner Galaxy has recently been re-investigated by Englmaier \\& Gerhard (1999). To keep this introduction as concise as possible, we quote here only the most relevant or recent papers and refer otherwise to the corresponding sections of MDZ96 and references therein. For the same reason we use a number of abbreviations whose meanings are explained in Table \\ref{abbtab}. For consistency, we adopt a distance to the Galactic Centre of $R_0 = 8.5$\\,kpc throughout this paper, although recent studies suggest a somewhat lower value (e.g., McNamara et al. 2000: $R_0 = 7.9\\pm 0.3$\\,kpc; see also Reid 1993).  MDZ96 classify the centre of our Galaxy as a mildly active Seyfert nucleus. Although the presence of a black hole of $\\sim 2.6\\times 10^6$\\,\\msun\\ is strongly supported by recent observations (e.g., Eckart \\& Genzel 1998) it has also become clear that most of the activity in the centre of our Galaxy is due to massive star formation in the central parsec (MDZ96, Sect. 5).  The mass of the central region of our Galaxy ($R \\le 3$\\,kpc) is dominated by the Galactic Bulge which consists mainly of old, evolved stars. Both the stellar near-infrared (NIR) surface brightness distribution and the kinematics of the gas suggest that the Galactic Bulge has a bar structure with its near end in the first galactic quadrant (i.e., at positive $l$) (e.g., Binney et al. 1991; Blitz \\& Spergel 1991; Weiland et al. 1994). As a consequence of the dynamics caused by the gravitational potential of this bar, the region around the co-rotation radius $R_{\\rm CR} \\sim (3.5\\pm 0.5)$\\,kpc is de-populated of gas, since Interstellar Matter (ISM) is transported efficiently inward from the {\\it Galactic Disk Molecular Ring} at the bar's outer Lindblad resonance at $R_{\\rm OLR} \\sim (4\\pm 0.5)$\\,kpc. Inside $R < 2$\\,kpc the gas settles on closed elongated ($X_1$) orbits. This gas is observed as a tilted disk of atomic hydrogen ($R \\le 1.5$\\,kpc), usually referred to as the {\\it ``H\\,I Central Disk''} (or {\\it ``Nuclear Disk''}) ($M_{\\rm H} \\sim 4\\times 10^7$\\,\\msun; Burton \\& Liszt, 1978). Inside the inner Lindblad resonance the $X_1$ orbits become self-intersecting, and shocks and angular momentum loss compress the gas into molecular form and drive it further inward where it finally settles on more circular, stable $X_2$\\ orbits (e.g., Englmaier \\& Gerhard 1999). The existence of a distinct, unusually dense molecular cloud complex in the central few hundred pc of our Galaxy, often referred to as the {\\it ``Central Molecular Zone''}, (CMZ) is well-established since the early 1970s (e.g., reviews by Genzel \\& Townes 1987 and MDZ96). The gas distribution in the CMZ is highly asymmetric with most of the mass being located at positive longitudes and positive velocities (see MDZ96, Sects. 2.2.3 and 3.4). One of the most remarkable features in the $l-v$\\ plane of the CMZ is the {\\it ``180-pc Molecular Ring''} which is hypothesized to be a shock region between the innermost stable $X_1$\\ orbit of the bar and the more circular $X_2$\\ orbits in the centre (Binney et al. 1991), but was also interpreted as an expanding molecular ring or shell (e.g., Scoville 1972; Bally et al. 1987; Sofue 1995b).  The large and dense stellar complex in this region was originally thought of as the innermost part of the more extended Galactic Bulge, with its population of old and evolved stars. Various observations, e.g. ongoing star formation, the presence of ionizing stars, and its extraordinary high surface brightness suggest that this region is distinct from the old Galactic Bulge (as originally proposed by Serabyn \\& Morris 1996) and may be associated with the CMZ. MDZ96 therefore call the innermost region $R < 300$\\,pc {\\it ``Nuclear Bulge''} (NB), and clearly distinguish it from the {\\it ``Galactic Bulge''} (GB). Due to its relative proximity, the physical characteristics of the NB of our Galaxy can be studied in detail.  However, the layer of interstellar dust in the Galactic plane restricts observations of the NB to wavelengths $\\lambda \\ge 2.2$\\,\\mim\\ and the edge-on projection makes a derivation of the true three-dimensional morphology difficult.  We have begun a systematic investigation of the physical characteristics of the NB the results of which are being published in a series of papers.  Papers\\,I (Philipp et al. 1999a) and II (Mezger et al. 1999) analysed the stellar population of the central $\\sim$30\\,pc. Based on a high-resolution $\\lambda$\\,2.2\\,\\mim\\  survey, we determined the K-band luminosity function (KLF) of the central 30\\,pc and interpreted it in terms of a present-day  bolometric luminosity and mass function. Main Sequence (MS) stars with masses $\\le 1$\\,\\msun\\ account for $\\sim$90\\% of the dynamical mass, but only for 6\\% of the K-band flux density. MS stars with masses $\\ge 1$\\,\\msun\\ account for $\\sim$6\\% of the dynamical mass and a similar percentage of the integrated K-band flux density, but are responsible for $\\sim$80\\% of the bolometric stellar luminosity as well as the ionization of the observed H{\\small II} regions. The bulk of the K-band emission comes from stars evolved from the MS such as giants, supergiants, and Wolf-Rayet stars. We find a deficiency of low-mass stars within the central 1.25\\,pc and indications of a high star formation activity during the past $10^7-10^8$\\,years.  Here, in paper III, we analyse the large-scale distribution of stars and ISM in the NB using IRAS and COBE DIRBE data. The size scales addressed in this paper span tens to hundreds of parsecs. The complex structure in the central few parsecs, including, e.g., the circumnuclear disk, the mini-spiral, and the central cavity, is not addressed. In a succeeding paper, based on ground-based single-dish mm observations, we investigate the morphology and kinematics of ISM in the central part of the NB in more detail (Zylka et al., in prep.).  The paper is organized as follows: Section \\ref{data} describes briefly the observational data used in this paper. Section \\ref{datan} describes the data reduction and analysis. In Sect. \\ref{results} we present the basic observational results, and in Sect. \\ref{disc} we derive physical characteristics of the NB and the CMZ. Section \\ref{nbmod} summarizes the results in terms of a coherent picture of the NB. ",
        "conclusions": ""
    },
    "0201/astro-ph0201541_arXiv.txt": {
        "abstract": " ",
        "introduction": "In recent years a standard model has emerged for the origin of the large-scale structure of our Universe~\\cite{LL}. Observations of the cosmic microwave background (CMB) reveal primordial anisotropies on the surface of last scattering of the CMB photons. Structure can form from these initial perturbations about a Friedmann-Robertson-Walker (FRW) universe by gravitational instability to form the galaxies, clusters of galaxies and superclusters observed in large-scale surveys. These observations are consistent with an almost scale-invariant initial spectrum of adiabatic density perturbations at last-scattering~\\cite{WTZ}. Inflation is a dynamical model of the early universe which can explain the origin of perturbations on arbitrarily large scales from small-scale vacuum fluctuations of light fields.  Observational data are increasingly being used to constrain the cosmological parameters of the background FRW model of the universe since last-scattering. But this can only be done in the context of some model for the nature of the primordial perturbations. Constraints on the form of the primordial perturbations yield constraints on the dynamics and high energy physics driving inflation in the very early universe. ",
        "conclusions": "It is quite possible that the primordial perturbations observed by future CMB experiments will remain consistent with the scale-invariant Gaussian spectrum of adiabatic density perturbations proposed more than thirty years ago by Harrison and Zel'dovich. In this case we would be able to extract little about the physics of inflation and would only be able to place bounds on allowed deviations from the extreme slow-roll limit.  If we are to learn more about the dynamical history of the early universe we need to find deviations from scale-invariance or traces of either tensor perturbations or non-adiabatic effects which, though not evident in current observations, could be detected by future experiments."
    },
    "0201/astro-ph0201407_arXiv.txt": {
        "abstract": "We have completed a new fiber array, SparsePak, optimized for low-surface-brightness studies of extended sources on the WIYN telescope. We are now using this array as a measuring engine of velocity and velocity-dispersion fields of stars and ionized gas in disk galaxies from high to low surface-brightness. Here we present commissioning data on the velocity ellipsoids, surface densities and mass-to-light ratios in two blue, high surface-brightness, yet small disks. If our preliminary results survive further observation and more sophisticated analysis, then NGC 3949 \\index{object, NGC 3949} has $\\sigma_z/\\sigma_R\\gg 1$, implying strong vertical heating, while NGC 3982's \\index{object, NGC 3982} disk is substantially sub-maximal. These galaxies are strikingly unlike the Milky Way, and yet would be seen more easily at high redshift.   \\vskip -0.25in ",
        "introduction": "\\vskip -0.05in  SparsePak contains a 76$\\times$77 arcsec sparsely-packed hexagonal grid of 75, 5-arcsec diameter fibers, which pipe light from the ``imaging'' Nasmyth focus to the versatile Bench Spectrograph on the WIYN 3.5m telescope (Bershady et al. 2002). Seven ``sky'' fibers are placed at $\\sim$ 30 arcsec from two grid sides. The large anamorphic demagnification in the echelle-mode of the Bench Spectrograph yields instrumental spectral resolutions of 8000-10,000 and as high as 21,000 with SparsePak -- comparable to resolutions obtained with smaller fibers. At the same time, SparsePak has $>90$\\% throughput for $\\lambda>500$nm, and provides nearly 3$\\times$ the light gathering power over 5$\\times$ the area as DensePak, the pre-existing WIYN fiber array. These attributes make SparsePak well suited for medium-resolution spectroscopy on extended objects down to low surface-brightness. ",
        "conclusions": ""
    },
    "0201/astro-ph0201517_arXiv.txt": {
        "abstract": "Recent observations of high redshift quasars at $z \\sim 6$ have finally revealed complete Gunn-Peterson absorption. However, this at best constrains the volume-weighted and mass-weighted neutral fractions to be $x_{\\rm HI}^{\\rm V} \\ge 10^{-3}$ and $x_{\\rm HI}^{\\rm M} \\ge 10^{-2}$ respectively; stronger constraints are not possible due to the high optical depth for hydrogen Lyman transitions. Here I suggest certain metal lines as tracers of the hydrogen neutral fraction. These lines should cause unsaturated absorption when the IGM is almost fully neutral, if it is polluted to metallicities ${Z \\sim 10^{-3.5}-10^{-2.5} Z_{\\odot}}$. Such a minimal level of metal pollution is inevitable in the middle to late stages of reionization unless quasars rather than stars are the dominant source of ionizing photons. The OI line at $1302$ \\AA$\\,$is particularly promising: the OI and H ionization potentials are almost identical, and OI should be in very tight charge exchange equilibrium with H. The SiII 1260 \\AA $\\,$transition might also be observable. At high redshift, overdense regions are the first to be polluted to high metallicity but the last to remain permanently ionized, due to the short recombination times. Such regions should produce a fluctuating OI and SiII forest which, if observed, would indicate large quantities of neutral hydrogen. The OI forest may already be detectable in the SDSS $z=6.28$ quasar. If seen in future high-redshift quasars, the OI and SiII forests will probe the topology of reionization and metal pollution in the early universe. If in additional the HI optical depth can be measured from the damping wing of a high-redshift gamma-ray burst, they will yield a very robust measure of the metallicity of the high-redshift universe. ",
        "introduction": "Recent spectroscopic observations \\cite{becker01,pent01,dvo01} of high-redshift quasars discovered by the Sloan Digital Sky Survey (Fan et al 2000, 2001a) have revealed long gaps in the spectra consistent with zero transmitted flux. This long-awaited detection of the Gunn-Peterson effect may herald the observational discovery of the reionization epoch. However, the high oscillator strength of the hydrogen Ly$\\alpha$ transition means that complete Gunn-Peterson absorption is expected even for a highly ionized intergalactic medium (IGM). The strongest constraint comes from the Ly$\\beta$ absorption trough, due the weaker (by $\\sim 5$) oscillator strength of Ly$\\beta$. From the $z=6.28$ quasar observed by \\scite{becker01}, \\scite{fan01b} conclude that at $z \\sim 6$, the lower limits on the mass-weighted and volume-weighted neutral hydrogen fraction are $x_{\\rm HI}^{\\rm M} > 10^{-2}$ and $x_{\\rm HI}^{\\rm V} > 10^{-3}$ respectively, larger by almost two orders of magnitude from $z \\sim 4$. Studies interpreting the observations conclude that the observed absorption troughs are consistent with the tail end of reionization, or post-overlap phase after individual HII regions have merged \\cite{barkana,fan01b}. However, due to the rapid or phase-change like nature of reionization in standard scenarios \\cite{gnedin,raz}, the ``dark ages'' or pre-overlap phase, when a substantial fraction of the hydrogen in the universe was neutral, is likely not far off. The spectra of the $z=6.28$ quasar suggests a very rapid evolution in the effective optical depth and thus the ionizing radiation field and effective neutral fraction \\cite{fan01b}. This implies that a slightly higher redshift quasar may indeed lie within the pre-overlap era.  Unfortunately, even if such a quasar is discovered, we may not learn anything new about the pre-reionization epoch. The hydrogen Lyman-series absorption trough saturates fully for a neutral hydrogen fraction at mean density $x_{\\rm HI} \\sim 10^{-4}$; because the transmitted flux declines exponentially with an increasing neutral fraction, we do not have the power to distinguish between an almost fully neutral IGM and one with only a tiny neutral fraction. As the IGM becomes almost fully neutral we might observe the red damping wing of the Gunn-Peterson trough \\cite{jordi}. Unfortunately, the highly luminous quasars presently observed probably ionize their surroundings on several Mpc scales; the consequent reduction in optical depth precludes observation of the red damping wing \\cite{cenhaiman,madaurees}. The only hope of detecting a damping wing would be to discover objects that ionize only a small region of the surrounding IGM: either a high-redshift gamma-ray burst (which has a very short duty cycle) or less luminous quasars or galaxies (which can be detected through gravitational lensing, \\pcite{ellis}). Alternatively, one might hope to detect the Ly$\\alpha$ halo surrounding a high-redshift source as Ly$\\alpha$ photons scatter and redshift in the surrounding neutral IGM \\cite{loebryb}. This also suffers from the difficulty that sources tend to ionize their surroundings; furthermore, the low surface brightness of the halo implies that detection is likely only possible with NGST.  What can be done with present-day technology? Clearly, we need absorption-line probes which are still unsaturated when the IGM is predominantly neutral. This is possible if the absorbers are much less abundant than hydrogen or have very small oscillator strengths. They should have ionization potentials similar to that of hydrogen in order to trace the HI fraction as faithfully as possible. In addition, their absorption lines must lie redward of the hydrogen Ly$\\alpha$ wavelength $\\lambda= 1216 \\, \\AA$, in order to avoid confusion with the lower-redshift Ly$\\alpha$ forest. In this paper, I suggest metal absorption lines as a probe of the neutral IGM. Metals are a natural probe: for a fully neutral IGM, $\\tau_{\\rm H \\, Ly\\alpha} \\sim 10^{5}$ at $z \\sim 6$, and while the oscillator strengths of metal UV/optical transitions ($f \\sim 10^{-2}-1$) are roughly comparable to that of hydrogen Ly$\\alpha$, the abundance by number of metals should be lower by $\\sim 10^{-6}-10^{-5}$, implying $\\tau_{\\rm metals} \\sim 10^{-2}-1$. The most uncertain aspect of this calculation is the degree to which an IGM polluted by metals can still remain neutral. I argue that because overdense regions are the first to be polluted with metals but the last to be permanently ionized (due to the short recombination time), a scenario of a neutral but metal-polluted IGM is plausible. Nonetheless, because of this uncertainty, a null detection of absorption will only yield a constraint on the joint metallicity/ionization state of the IGM. A positive detection, however, may be our best hope of unveiling an almost fully neutral IGM with observations of high-redshift quasars in the near future. In all numerical estimates, I assume a $\\Lambda$CDM cosmology with $(\\Omega_M,\\Omega_\\Lambda,\\Omega_b,h, \\sigma_{8 h^{-1}},n)=(0.35,0.65,0.04,0.65,0.87,0.96)$. ",
        "conclusions": "The SDSS 1030+0524 quasar at $z=6.28$ shows tantalizing absorption features blueward of the OI 1305${\\rm \\AA}$ line \\cite{becker01}. There also appears to be a fairly deep absorption feature blueward of the SiII 1260$\\AA$. Could the absorption lines described in this paper been seen already? Unfortunately, the quasars SDSS 1044-0125($z=5.80$), 0836+0054($z=5.82$),1306+0356($z=5.99$) also show some absorption features in the same wavelength interval; in particular, SDSS 1306+0356 shows a very strong absorption feature at $\\sim 7130 {\\rm \\AA}$, with no detected flux over $\\sim 80 {\\rm \\AA}$ (this has been tentatively identified as CIV absorption at $z=4.86$). These features cannot correspond to OI absorption lines: unless they correspond to regions of anomalously high metallicity, the associated hydrogen column densities would be $N_{HI} > 10^{20} {\\rm cm^{-2}}$, and all flux at the hydrogen Ly$\\alpha$ wavelength should be obliterated, while some flux is still seen there. These lines are probably associated with metal lines (e.g. MgII) from lower redshift absorbers, and illustrate a generic difficulty in observing the features proposed in this paper. On the other hand, OI lines are not ruled out in the $z=6.28$ quasar because of the complete damping of flux at Ly$\\alpha$, Ly$\\beta$ wavelengths. As we have seen, these absorption features can still arise in the post-overlap epoch when regions with $\\Delta > {\\rm few}$ are largely neutral; indeed, up to a few absorption lines with observed equivalent widths $W_{\\lambda} \\sim 5 \\redshift \\left( \\frac{N_{i}}{10^{15} \\, {\\rm cm^{-2}}} \\right)$ \\AA $\\,$ might be seen. Note that the spectra of \\scite{becker01} have been smoothed to 4\\AA $\\, {\\rm pixel}^{-1}$. The absorption features blueward of OI 1305${\\rm \\AA}$ cannot be definitely identified as OI absorption. Two of them lie at the same wavelengths as bright sky emission lines and are probably due to imperfect sky subtraction. A third line with observed frame equivalent with $\\sim 25 \\, \\AA$ is a possible candidate; however, it lies very close to the rest frame OI wavelength and could be due to intrinsic absorption. More cannot be said without careful study of the spectrum. A definitive detection of the OI forest can probably only be done with much higher signal-to-noise spectra of the same quasar.  The estimates in this paper can be addressed with in much greater detail with numerical simulations. In particular, I used very simple ansatzes for the dependences of metallicity and ionization fraction with overdensity, $Z(\\Delta), x_{i}(\\Delta)$ which in fact should be highly stochastic and spatially varying. They can be much better modelled in a self-consistent fashion in simulations which attempt to model the metal pollution \\cite{cenostriker,anthony} and radiative transfer \\cite{gnedin,raz}, particularly since the rise in metallicity and ionization fraction are inter-related. The spatial structure of the OI forest can be computed by shooting lines of sight through a simulation box. If the OI forest is indeed seen, such studies will be urgently needed to provide a more realistic interpretation of the observations.  The scenario in this paper is not significantly altered if a substantial X-ray background due to high redshift supernovae \\cite{oh2001} or quasars \\cite{venkatesan} is present. The large mean free path of hard photons means that they can ionize the IGM fairly uniformly, but beyond $x_{e} \\sim 0.1$ most of the energy of an energetic electron created goes into Coulomb heating the gas rather than collisional ionization \\cite{shull}; predominantly neutral regions should therefore still exist.  We will gain a wealth of information about early metal pollution and the reionization process if the OI and SiII forests are seen. They will be direct probes of the topology and history of gas clumping, metal pollution and reionization in the early universe. As we have seen, if we assume a relation between the metals and ionizing photons produced by massive stars, they could potentially also provide indirect constraints on the escape fraction of ionizing photons from star forming halos and the QSO contribution to the ionizing background, and the filling factor of metal pollution. They may also give clues as to the nature of the ionizing sources: the structure of forest lines should look different if the universe were reionized by rare but luminous source as opposed to abundant but faint sources, since in the former case higher overdensity regions have to be ionized before overlap can be achieved. OI and HI will be locked in very tight charge exchange equilibrium at high redshift. If we are lucky enough to observe the rest-frame optical afterglow of a high-redshift gamma-ray burst and measure both $\\tau_{\\rm HI}^{\\rm eff}$ (from the damping wind) and $\\tau_{\\rm OI}^{\\rm eff}$, we will have a direct measure of the mean metallicity of the universe at high redshift, independent of gas clumping or the form of the ionizing radiation field. Otherwise, a lower limit on the metallicity from measurement of $\\tau_{\\rm OI}^{\\rm eff}$ alone is possible. A null detection of the OI, SiII forests will yield constraints on the parameter $Y_{i}(\\Delta)= x_{i}(\\Delta) Z(\\Delta)$, but a positive detection will be tremendously exciting and almost certainly signal the presence of almost fully neutral hydrogen at high redshift. To date, the OI and SiII forests may be our only probes of nearly neutral gas in the pre-reionization epoch observable with current technology."
    },
    "0201/astro-ph0201384_arXiv.txt": {
        "abstract": "We report the discovery with the Parkes radio telescope of a pulsar associated with the $\\sim 1700$\\,yr-old oxygen-rich composite supernova remnant \\snr.  \\psr\\ has period 135\\,ms and period derivative $7.4\\times10^{-13}$, implying characteristic age 2900\\,yr, spin-down luminosity $1.2\\times10^{37}$\\,erg\\,s$^{-1}$, and surface magnetic field strength $1.0\\times10^{13}$\\,G.  Association between the pulsar and the synchrotron nebula previously identified with \\chandra\\ within this supernova remnant is confirmed by the subsequent detection of X-ray pulsations by Hughes et~al.  The pulsar's flux density at 1400\\,MHz is very small, $S \\approx 80\\,\\mu$Jy, but the radio luminosity of $S d^2 \\sim 2$\\,mJy\\,kpc$^2$ is not especially so, although it is one order of magnitude smaller than that of the least luminous young pulsar previously known.  This discovery suggests that very deep radio searches should be done for pulsations from pulsar wind nebulae in which the central pulsed source is yet to be detected and possibly other more exotic neutron stars. ",
        "introduction": "\\label{sec:intro}  The supernova remnant (SNR) \\snr\\ is one of only three oxygen-rich SNRs known in the Galaxy.  The other two (Puppis~A and Cas~A) have central compact objects, the nature of which however remains mysterious (e.g.\\ \\cite{gbs00}; \\cite{pza+00}). At radio wavelengths, \\snr\\ has the appearance of a composite SNR, with a central peak and a shell $\\sim 10'$ in diameter (\\cite{bgcr86}).  ASCA X-ray observations (\\cite{tts98}) detected a non-thermal nebula coincident with the central radio component.  Recent \\chandra\\ (ACIS-S) observations have shown this nebula to be $\\sim 2'$ in extent, and to contain a resolved compact source located near its peak.  This discovery, together with the energetics of the nebula, provides nearly incontrovertible evidence for the existence of a pulsar powering the nebula (\\cite{hsb+01}).  The \\chandra\\ observations of \\snr\\ are a beautiful example of the recent flood of X-ray data with high spatial, temporal and spectral resolution that is advancing dramatically our understanding of the varied outcomes of supernova explosions.  In particular, X-rays provide an important complement to the radio band, the traditional hunting ground of pulsar studies.  In this Letter we report the discovery and key parameters of the pulsar in \\snr\\ in a deep observation with the Parkes radio telescope.  The characterization of this young and energetic pulsar is important for the analysis of existing X-ray and radio data on \\snr.  Moreover, this discovery has significant implications for the concept of ``radio-quiet neutron stars.'' ",
        "conclusions": "\\label{sec:disc}  The measured $P$ and $\\dot P$ of \\psr\\ imply that it is a very young and energetic pulsar.  The implied characteristic age and surface magnetic dipole field strength are $\\tau_c = P/2\\dot P = 2900$\\,yr and $B = 3.2\\times10^{19} (P \\dot P)^{1/2} = 1.0 \\times 10^{13}$\\,G respectively.  The spin-down energy loss rate $\\dot E = 4\\pi^2 I \\dot P /P^3 = 1.2 \\times10^{37}$\\,erg\\,s$^{-1}$ (where a neutron star moment of inertia $I = 10^{45}$\\,g\\,cm$^2$ has been used), in excellent agreement with the value predicted by Hughes et al. (2001)\\nocite{hsb+01} from the energetics of the PWN, if $d \\sim 5$\\,kpc.  In comparison with the sample of $\\sim 1500$ rotation-powered pulsars known, J1124$-$5916 ranks as the 6th youngest in terms of $\\tau_c$ and the 8th most energetic in terms of $\\dot E$.  The age of \\snr\\ is derived to be $\\la 1700\\,(d/5\\,\\mbox{kpc})$\\,yr from a measurement of the high radial velocity oxygen-rich material positionally coincident with the central synchrotron nebula (\\cite{mc79}; \\cite{bgdb83}), while the pulsar characteristic age is a factor of nearly two larger than this.  Assuming that the age estimate for the SNR represents the true age of the system, this discrepancy can be simply explained if the initial spin period of \\psr\\ was $P_0 \\ga 90$\\,ms (e.g.\\ \\cite{krv+01}; \\cite{mss+01}).  This is slower than for the six other young pulsars whose initial periods have been estimated, all of which have $P_0 < 60$\\,ms.  Alternatively the pulsar may spin down with a braking torque that is different from the usually assumed constant magnetic dipole.  In a study of pulsar population dynamics, Cordes \\& Chernoff (1998)\\nocite{cc98} suggest alternate forms of braking torque (e.g.\\ a braking index of 2.5 with a magnetic field decay time $\\sim 6$\\,Myr) that result in actual pulsar ages smaller than $\\tau_c$ by a factor $\\la 2$.  In the absence of a measured braking index for \\psr, we tentatively regard a relatively large $P_0$ as a preferred explanation.  In Table~\\ref{tab:young} we summarize key parameters for the youngest rotation-powered pulsars known, ordered by decreasing $\\dot E$. Heading the list are the three ``Crab-like'' pulsars, with extremely large $\\dot E$.  The remaining six young pulsars are varied in their properties and generally follow the trend of increasing $P$ and $B$ with decreasing $\\dot E$ (with the pulsar in G11.2$-$0.3 a notable exception).  Their spin-down luminosities are 20--200 times lower than the Crab's and many are therefore difficult to detect: four were discovered just in the past two years (two each in radio and X-rays).   In its spin parameters \\psr\\ is most similar to PSR~B1509$-$58 (see Table~\\ref{tab:young}).  It is interesting to note that their respective PWNe also have comparable luminosities.  In the 0.2--4\\,keV X-ray band, we find for the PWN powered by J1124$-$5916 that\\footnote{The photon index of this source could not be measured from the \\chandra\\ observations of Hughes et al. (2001\\nocite{hsb+01}) because of contamination by thermal emission from the SNR.  We assume a photon index $\\Gamma = 2$, as is typical for such sources.} $L_X = 6\\times10^{34}$\\,erg\\,s$^{-1} = 0.005 \\dot{E}$ for a distance 5\\,kpc (\\cite{hsb+01}), while for B1509$-$58 we find $L_X = 2\\times10^{35}$\\,erg\\,s$^{-1} = 0.01 \\dot{E}$ in the same energy range (\\cite{gak+01}).  For their respective radio PWNe, we find $L_R = 4\\times10^{33}$\\,erg\\,s$^{-1} = 0.0003 \\dot{E}$ for J1124$-$5916 (\\cite{hsb+01}), while $L_R \\sim 5\\times10^{33}$\\,erg\\,s$^{-1} = 0.0003 \\dot{E}$ for B1509$-$58 (\\cite{gak+01}).  Thus the efficiencies with which these pulsars power their X-ray nebulae differ only by a factor of $\\sim2$, while those for their radio nebulae are approximately equal.  Although these pulsars have comparable spin parameters, we explain below that their PWNe have quite different environments and likely evolutionary histories. It is therefore surprising that the efficiencies with which they convert their spin-down into nebular flux are so similar.  In the radio band, it has been argued that a pulsar's initial parameters can make a crucial difference to the resultant luminosity of its PWN at later times (e.g.\\ \\cite{bha90}).  If a pulsar is born spinning rapidly and has a high magnetic field, it will dump most of its spin-down energy into its surrounding nebula at the earliest stages, when adiabatic losses are most severe; this will result in a comparatively underluminous radio PWN (\\cite{bha90}; \\cite{cgk+01}). However, a pulsar that is otherwise similar but born spinning more slowly will not release its energy so quickly, so that its nebula will suffer less from expansion losses and will be correspondingly brighter at radio wavelengths. We have argued above that \\psr\\ was possibly born with a comparatively long spin period, $P_0 \\ga 90$\\,ms, while PSR~B1509$-$58 is generally assumed to have been born with a Crab-like $P_0 \\approx 20$\\,ms (\\cite{bha90}).  The comparable values of $L_R$ for their PWNe are therefore not easily explained.  Furthermore, the fact that the radio and X-ray extents of the PWN powered by PSR~B1509$-$58 are approximately equal indicates that synchrotron cooling does not yet dominate the nebula at high energies (\\cite{gak+01}), and can account for this nebula's comparatively low X-ray efficiency (\\cite{che00}). However, for \\psr\\ the X-ray nebula is noticeably smaller than its radio counterpart (\\cite{hsb+01}; Gaensler \\& Wallace, in preparation), indicating that the X-ray emitting electrons are in this case efficient radiators.  Thus we expect the X-ray luminosity of the PWN around J1124$-$5916 to be much larger than observed, closer to the factor $L_X \\approx 0.05 \\dot{E}$ seen for the Crab Nebula and other X-ray PWNe dominated by radiative losses (\\cite{che00}).  Thus both $L_R$ and $L_X$ for the PWN powered by \\psr\\ are much lower than expected through simple physical arguments.  There are a variety of other factors which can affect a PWN's luminosity: the pulsar braking index, PWN magnetic field, and nebular expansion velocity all are important parameters at radio wavelengths (\\cite{rc84}; \\cite{cgk+01}), while the Lorentz factor of the wind, the radius of the termination shock and the magnetization parameter, $\\sigma$, all have a strong bearing on the nebula's X-ray luminosity (\\cite{che00}).  The unexpectedly similar nebular efficiencies for PSRs~J1124$-$5916 and B1509$-$58, along with the large range of efficiencies among the other young pulsars in Table~\\ref{tab:young}, emphasize that we still lack a detailed understanding of how a pulsar ultimately deposits its energy.  One particularly interesting young pulsar in this regard is PSR~J1119$-$6127 (\\cite{ckl+00}; Table~\\ref{tab:young}), which has no known radio-bright PWN down to extremely constraining surface brightness limits (\\cite{cgk+01}).  It also has no confirmed X-ray-bright PWN (although a possible candidate was detected by \\cite{pkc+01}).  This example, together with several relatively young and energetic pulsars that do not power detectable PWNe (\\cite{gsf+00}), suggests that pulsar environments need not be good ``calorimeters.''  Thus, some young pulsars may pass unnoticed unless their beamed radiation is favorably oriented so that pulsations may be detected.  That radio pulsar beams do not in general sweep $4\\pi$\\,sr is a fact of life.  Although it is generally accepted that the beaming fraction $f$ is period dependent, with short-period pulsars beaming to a larger fraction of the sky than their long-period counterparts, a consensus has yet to emerge on the form of $f(P)$.  Pulse width analyses suggest $f \\approx 0.3$ for a pulsar with $P \\sim 0.1$\\,s (\\cite{tm98}) while, from an analysis of pulsar--PWN associations, Frail \\& Moffett (1993)\\nocite{fm93} find $f \\approx 0.6$. Since PWNe are regarded as unambiguous indicators that a young pulsar is present, the failure to detect pulsations towards such sources is usually ascribed solely to a misoriented pulsar beam.  However, the case of \\snr\\ and its faint radio pulsar highlights an important caveat: the radio luminosity\\footnote{The radio luminosities we discuss here are really ``pseudo-luminosities'': they assume that total luminosity is proportional to the integrated flux density of the observed cut across the radio beam.  A realistic discussion of actual luminosities depends crucially on the generally unknown pulsar beam shape: e.g., for a beam with exponential roll off, the measured flux density --- and therefore pseudo-luminosity --- can be arbitrarily low depending on the viewing impact angle, even for a large beam-averaged luminosity.  Here primarily we are concerned with empirically determined ``luminosities,'' i.e., with pseudo-luminosities.} limit implied by a non-detection should now be below at least $L_{1400} \\sim 1$\\,mJy\\,kpc$^2$, approximately the luminosity of \\psr, and possibly as low as $\\la 0.1$\\,mJy\\,kpc$^2$ (see Fig.~\\ref{fig:Lr}), before one can invoke unfavorable beaming as an explanation.  Various analyses (e.g.\\ \\cite{ec89}) have suggested that young pulsars have higher radio luminosities than older ones.  Figure~\\ref{fig:Lr} shows that the basis for such arguments is weak in the case of very young pulsars.  The least luminous radio pulsar previously detected for which $\\tau_c < 10$\\,kyr had $L_{1400} \\sim 30$\\,mJy\\,kpc$^2$.  The fact that some X-ray-detected pulsars have had much more constraining luminosity limits put on their radio emission (e.g.\\ \\cite{ckm+98}; Table~\\ref{tab:young}) has usually been taken to imply that these pulsars have radio beams which are not directed towards us.  However, the discovery of \\psr\\ shows that very young pulsars can certainly have much lower radio luminosities than previously thought, either due to intrinsically low luminosity or to our line of sight cutting only low-level wings on the emitted beam.  The above discussion serves as a cautionary tale in the context of discussions of exotic sources such as the ``radio-quiet neutron stars'' (RQNSs), ``anomalous X-ray pulsars'' (AXPs) and ``soft $\\gamma$-ray repeaters'' (SGRs). While there is no question that these sources have properties which make them distinct from ``normal'' radio pulsars (e.g.\\ \\cite{gpv99}; \\cite{pkc00}; Gaensler et al.~2000a\\nocite{gbs00}; \\cite{kgc+01}; \\cite{cph+01}), the fact that none of these sources has been detected at radio wavelengths has been taken to imply that their radio pulse mechanism must be inactive (\\cite{bh98b}; \\cite{zh00}; \\cite{zha01}).  However, the radio luminosity limits for many of these sources are not particularly stringent by the standard of radio pulsars, and there are so few of these objects known that they all could be radio luminous and yet all could still be beaming away from us (\\cite{gsgv01}).  It may well be that the RQNSs, AXPs and SGRs are all truly radio-silent, but this conclusion is difficult to justify until the luminosity limits available for such objects reach a level equivalent to at least the smallest radio luminosities observed for radio pulsars, $L_{1400} \\sim 0.1$\\,mJy\\,kpc$^2$ (Fig.~\\ref{fig:Lr}), and in principle even smaller values.  With the discovery of J1124$-$5916, \\snr\\ becomes the second example of an oxygen-rich SNR with a confirmed rapidly spinning pulsar (after 0540$-$693 in the LMC).  It remains to be seen what is the true nature of the compact sources in the other two Galactic oxygen-rich SNRs. This discovery also suggests that very deep searches for radio pulsations from known PWNe and possibly other more exotic neutron stars may be well worthwhile."
    },
    "0201/astro-ph0201451_arXiv.txt": {
        "abstract": "{ We present high resolution blue spectroscopy of an almost complete sample of optical counterparts to massive X-ray binaries in the Large Magellanic Cloud (LMC) and derive their spectral classification. We find an spectral type B0II for the optical counterpart to \\object{RX J0532.5$-$6551}, confirming it as the first wind-fed massive X-ray binary in the LMC. We also confirm the Be nature of the proposed counterpart to \\object{RX J0535.0$-$6700}. The proposed optical counterpart to  \\object{RX J0531.5$-$6518} is a B2V star with signs of emission in the Balmer lines. In total, we give accurate spectral types for 14 counterparts. We find that the overall observed population of massive X-ray binaries in the LMC has a distribution not very different from the observed Galactic population and we discuss different selection effects affecting our knowledge of this population. The spectral distribution of the Be/X-ray binary population is also rather similar to the Galactic one. This distribution implies that Be/X-ray binaries must have preferentially formed from moderately massive binaries undergoing semi-conservative evolution. The observation of several Be/X-ray binaries with large eccentricities implies then the existence of supernova kicks. } ",
        "introduction": "High Mass X-ray Binaries (HMXBs) are X-ray sources composed of an early-type massive star and an accreting compact object (generally a neutron star, but occasionally a black hole). HMXBs are traditionally divided (see Corbet 1986) into Classical or Supergiant X-ray  binaries (SXBs) in which the compact object accretes from  a mass-losing OB supergiant or bright giant and Be/X-ray binaries (BeXBs), in which a neutron star orbits an unevolved OB star surrounded by a dense equatorial disc (cf. Liu et al. 2000 for a recent catalogue), though the physical reality could be rather more complex (cf. Negueruela \\& Reig 2001). Different population synthesis analyses predict that the vast majority of HMXBs will be BeXBs, though this is not apparent from the number of sources detected in the Milky Way (where $\\sim 30$\\% of known systems are SXBs), presumably due to different selection effects that will be detailed later on.  Apart from their intrinsic interest as high-energy radiation sources, HMXBs offer a window on the late stages of massive binary evolution. Their properties as a population can be used to extract valuable information on the different stages of the life of massive binaries.  In order to understand the representativity of the known sample of Galactic HMXBs, it is of fundamental importance to have at least an approximate idea of the HMXB population content in other Galaxies. The Magellanic Clouds (MCs) present a unique opportunity to carry out such a study, since they have a structure and chemical composition which differs from that of the Milky Way and, at the same time, are close enough to allow the study of individual sources with modest-sized ground-based telescopes.  For this reason, we have undertaken an observational campaign in order to obtain high Signal-to-Noise Ratio (SNR) spectroscopy of the optical counterparts to MC HMXBs which will allow us to derive accurate spectral classifications for these stars. Such work is imperative if we are to gain some understanding of the mass distribution and evolutionary status of the HMXB populations.  In this first paper, we centre on the HMXB population of the LMC. This population is relatively small and all the optical counterparts are reasonably bright, allowing us to obtain a basically complete sample. In further works, we will study the SMC population and will compare the characteristics of the MC populations with the Galactic one. ",
        "conclusions": "\\label{sec:dis}  We have obtained accurate spectral classifications for an almost complete sample of HMXBs in the LMC. Among these objects, there are three well-known bright persistent X-ray sources, two of which are considered to be black hole candidates (LMC X-1, O8III-V; LMC X-3, B2?Vp), while the third is a bright X-ray pulsar fed by Roche-Lobe overflow from the O8III companion (LMC X-4). In addition, we have found that the optical counterpart to \\object{RX J0532.5$-$6551} has a spectral type B0II and therefore this source is likely to be a wind-fed accreting system, like Vela X-1. All the other sources in the sample are likely Be/X-ray binaries. Eight of them had been considered as such in previous work (of which, 5 are X-ray pulsars) and two further suggestions seem to be confirmed by our results. The list of confirmed and likely Be/X-ray binaries in the LMC is summarized in Table~\\ref{tab:bex}, while the list of possible HMXBs which have not been included in this investigation is given in Table~\\ref{tab:unknowns}.  Since we have relied on the derived absolute magnitudes of our objects to determine their luminosity classes, the choice of a particular calibration has a direct bearing on the results. Throughout Section~\\ref{sec:spec}, we have used the calibration of Vacca et al. (1996) for O-type stars, complemented by that of Schmidt-Kaler (1982) for early B stars. By using this calibrations, we have found that three objects seem to be underluminous for their spectral type: \\object{LMC X-4} is moderately underluminous for O8V, though a spectral type is difficult to assign. On the other hand, both \\object{1A\\,0538$-$66} and \\object{H0544$-$665} have well defined spectral types and their absolute magnitudes seem to be slightly lower than expected.  The distance modulus to the LMC used through this work is $(M-m)_{0}=18.24$, after Udalsky (2000). The validity of this value is still open to discussion, with many authors supporting a value $(M-m)_{0}=18.5$. We just note here that the distance modulus adopted does not affect in any significant way the results obtained. The only changes that would result from adopting the ``longer'' distance would be that the intrinsic luminosity of the optical companion of \\object{LMC X-4} would become marginally compatible with an O8V spectral type (but still not with O8III) and that the intrinsic luminosity of \\object{EQ 050246.6$-$663032.4} (CAL E) would then be compatible with a B0III spectral type. \\object{H0544$-$665} and \\object{1A\\,0538$-$66} would still be underluminous for their spectral types.  A new  calibration based on Hipparcos distances has been derived by Wegner (2000), who finds on average rather lower absolute magnitudes at a given spectral type than previously assumed (typically by $\\sim 1$ mag). By using this calibration, we could solve the problem of our three underluminous objects, but we would create a larger one, since the magnitudes that we derive (typically $M_{V}\\approx -4$ for B0V objects) would then mean that all the other counterparts are bright giants. Therefore we conclude that our data support the older ``brighter'' calibrations, even though we have used a ``short'' distance for the LMC.  \\subsection{A preliminary comparison to the Galactic population}  Our sample of the LMC HMXB population is complete, in the sense that we have observed all the systems with confirmed optical counterparts. As shown in Table~\\ref{tab:unknowns}, only two possible Be/X-ray binaries are left out (\\object{RX J0516.0$-$6916} and \\object{RX J0532.4$-$6535}). The main limitation to any conclusions to be drawn from this sample (i.e., how representative it is?) stems from our inability to judge how many other X-ray sources remain undiscovered. In this respect, it is clear that, unlike in the Galaxy, we are not missing any sources because they are very absorbed. The {\\em ROSAT} PSPC pointed at most locations in the LMC more than once, with many fields being observed close to ten times \\cite{hap99}. Since the sensitivity limit of {\\em ROSAT} allowed the detection of relatively weak LMC sources, only very weak persistent X-ray sources (with $L_{\\rm x}\\la10^{34}\\:{\\rm erg}\\,{\\rm s}^{-1}$) may have been missed, specially in fields where no deep exposures have been carried out.  The main uncertainty is therefore what fraction of transient sources has not been discovered yet because they were not active when they were looked at. There is no obvious way to assess this number, but a comparison with the Galactic sources may indicate that the census is still relatively incomplete, since a large fraction of Be/X-ray transients are only active as X-ray sources for relatively short periods.  At first sight, the LMC HXMB population is not significantly different from the Galactic one. All the LMC systems seem to have Galactic equivalents. \\object{LMC X-4} is similar to the Roche-Lobe overflow accreting pulsar \\object{Cen X-3}. \\object{LMC X-1} would be its equivalent with a black hole companion, not too different from \\object{Cyg X-1}, except in the mass transfer mechanism. \\object{LMC X-3} is similar to the recently found black hole candidate \\object{SAX J1819.3$-$2525}/\\object{V4641 Sgr} for whose components (B9III+BH) Orosz et al. (2001) derive masses which are quite close to those estimated for \\object{LMC X-3}. The fact that \\object{SAX J1819.3$-$2525} is a transient X-ray source, while \\object{LMC X-3} is persistent, may be related to the smaller orbit of \\object{LMC X-3}.  Among the list of known and probable HMXBs of Liu et al. (2000) we find that there are 40 Galactic systems with identified optical counterparts -- leaving aside some objects like \\object{$\\gamma$ Cas}, \\object{1E\\,1024.0$-$5732} or \\object{1H\\,0521+373} whose nature as HMXBs is unclear. Of these, 23 are BeXBs and 10 are SXBs. 7 objects are not included in any of the two groups. In particular, we exclude the binaries \\object{1E\\,0236.6+6100} and \\object{SAX J0635+0533} from the count of Be/X-ray binaries, because they are not believed to be accretion driven. We have also excluded \\object{XTE J0421+560}, because the exact nature of its two components is still unknown. The respective fractions are then 25\\% SXBs and 58\\% BeXBs. Adding objects without optical counterparts which are believed to belong to one of the two groups because of their X-ray properties or orbital solutions gives 13 SXBs and 37 BeXBs. The fraction SXBs/BeXBs among systems with identified counterparts is 0.43, which reduces to 0.35 when candidate unidentified sources are added.  Of the 10 Galactic SXBs (and 13 likely SXBs), only one (\\object{Cen X-3}) is powered by Roche-lobe overflow (RLO). All the others are wind-fed systems. Moreover, three of the Galactic HMXBs that have not been included in any of the two groups are also wind-fed X-ray sources: \\object{RX J1826.2$-$1450} and \\object{4U\\,2206+54} contain main-sequence O-type stars (c.f. Negueruela \\& Reig 2001) and \\object{Cyg X-3} likely contains a Wolf-Rayet star. This means that we know of 16 wind-fed X-ray sources in the Galaxy (most of them persistent and with $L_{\\rm x}\\ga10^{35}\\:{\\rm erg}\\,{\\rm s}^{-1}$) against 1 (or maybe 2 if \\object{RX J0541.4$-$6936} really belongs to this category) such sources in the LMC. Since some wind-fed systems are weaker X-ray sources ($L_{\\rm x}\\approx10^{34}\\:{\\rm erg}\\,{\\rm s}^{-1}$), the Galactic sample may not be complete. A simple mass ratio argument, if the stellar populations of both Galaxies are similar, would predict a ratio of 10\\,-\\,12 \\cite{sun92} more X-ray binaries in the Milky Way than in the LMC. Therefore the low number of wind-fed systems in the LMC is not surprising in principle.  What sets out a difference is the fact that for one or two wind-fed supergiants, there are 3 RLO persistent bright sources in the LMC. This is in contrast to the Galaxy, where only one such system (two RLO sources with massive companions, if the transient black hole candidate \\object{SAX J1819.3$-$2525} is counted) is known. Moreover, the three donors in the LMC RLO sources are quite close to the main sequence, while the donor in \\object{Cen X-3}, as the donor in \\object{SMC X-1} -- the only similar source in the Small Magellanic Cloud --, is rather more evolved. Since no selection effects are expected to be so strong as to affect the number of persistent sources with $L_{\\rm x}\\approx10^{38}\\:{\\rm erg}\\,{\\rm s}^{-1}$ detected in our Galaxy, and a mass ratio argument would predict a much larger number of RLO objects, there is a strong suggestion here of some significant difference. Though the numbers involved are too small to attempt any statistics, there is good reason to suspect that, for some reason, binary evolution is more likely to result in close binaries with black holes in the LMC than in the Milky Way, as has been discussed by several authors (e.g., Johnston et al. 1979; Pakull 1898). This seems to be the only significant difference among the two populations and it is unclear whether it can be assigned to the effects of lower metallicity in binary evolution (e.g., through weaker stellar winds resulting in higher pre-supernova core masses) -- see Helfand \\& Moran (2001) for a more thorough discussion.  Among the objects in our LMC sample, the fraction of BeXBs is 71\\%. If all the proposed HMXBs in Table~\\ref{tab:unknowns} are added, the proportion is 70\\%. The corresponding fraction for the Galactic sample (i.e. BeXB/HMXB) is 58\\% for systems with optical counterparts and 65\\% for all systems (i.e., identified + candidates). The fact that the fraction is so similar in both galaxies strongly suggests that the selection effects dominating our knowledge of the HMXB population are similar -- i.e., in both cases, the main bias results from the incompleteness of the Be/X-ray binary sample, due to their transient X-ray source condition.  \\begin{table}[ht] \\caption{ Be/X-ray binaries in the LMC with their spectral types and other known parameters. See the text for references.} \\begin{center} \\begin{tabular}{lccc} \\hline Name & $P_{{\\rm s}}$(s) & Spectral & Max $L_\\mathrm{x}$\\\\ & & Type &  (erg s$^{-1}$)\\\\ \\hline \\object{CAL 9} & $-$ & B0V & $7\\times10^{34}$\\\\ \\object{CAL E} & 4.1 & B0V& $4\\times10^{37}$\\\\ \\object{RX J0520.5$-$6932} & $-$ & O9V & $5\\times10^{34}$ \\\\ \\object{RX J0529.8$-$6556} & 69.5 & B0.5V& $2\\times10^{36}$ \\\\ \\object{EXO 0531.1$-$6609} & 13.7 & B0.7V& $1\\times10^{37}$ \\\\ \\object{RX J0531.5$-$6518} & $-$ & B2V &$3\\times10^{35}$ \\\\ \\object{RX J0535.0$-$6700} & $-$ &B0V& $3\\times10^{35}$ \\\\ \\object{1A\\,0535$-$66} & 0.07 & B0.5III& $1\\times10^{39}$\\\\ \\object{1SAX J0544.1$-$7100} & 96.1 & B0V & $2\\times10^{36}$\\\\ \\object{H0544$-$665} & $-$ & B0V & $1\\times10^{37}$\\\\ \\hline \\end{tabular} \\end{center} \\label{tab:bex} \\end{table}  \\subsection{The spectral distribution of Be/X-ray binaries}  The spectral distribution  of the optical counterparts to Galactic Be/X-ray binaries was studied by Negueruela (1998), who found it to be very different from the spectral distribution of isolated Be stars. The spectral distribution of optical counterparts peaks sharply at B0 and does not extend beyond B2 (roughly corresponding to $\\approx10\\,M_{\\sun}$). Such distribution can be explained if one assumes that during the mass transfer phase previous to the formation of the Be/X-ray binary (see next Section) material lost from the system carries away a large amount of angular momentum \\cite{vbv97}.  \\begin{figure} \\begin{picture}(250,360) \\put(0,0){\\special{psfile=MS2114f10.ps angle =0 hoffset=-20 voffset=-5 hscale=48 vscale=48}} \\end{picture} \\caption{The spectral distribution of the optical counterparts to Be/X-ray binaries in the LMC (filled) is compared to the distribution of 16 counterparts to Galactic Be/X-ray binaries identified with X-ray pulsars (hollow). Negative spectral subtypes are used to represent O-type stars.} \\label{fig:hist} \\end{figure}  The spectral distribution of optical counterparts to LMC Be/X-ray binaries is displayed in Fig.~\\ref{fig:hist}, together with that of Galactic sources. The convention adopted has been that of Steele et al. (1998), i.e., the B0 bin contains objects with spectral types in the range O9.5 to B0.2. In order to allow better comparison, the Galactic sample has been replotted using the same convention. This sample contains the 13 X-ray pulsars listed in Negueruela (1998) and three new pulsars discovered by Motch et al. (1997): \\object{RX J0440.9+4431}, \\object{RX J0812.4$-$3114} and \\object{RX J1037.5$-$5647}.  A detailed statistical comparison is left for future work, after the spectral types of some Galactic objects have been reassessed in view of new high-quality spectra. However, it is clear from Fig.~\\ref{fig:hist} that the two distributions are basically identical. Though such similarity is in principle expected (see Van Bever \\& Vanbeveren 1997), it is encouraging to see it confirmed, because it must mean that: \\begin{itemize} \\item The Galactic sample, in spite of all the possible selection effects (see discussion in Negueruela 1998) is representative of the actual spectral distribution. \\item In spite of its small size (which could make us doubt of its statistical significance), the LMC sample represents well the population. \\end{itemize}  If selection effects are having any impact on the spectral distribution observed, such effects must affect both samples. However, selection effects should be very different in the Galactic plane (where we expect extinction towards the optical companions to be the main source of bias in our knowledge of the spectral distribution -- in the sense that earlier spectral types are easier to observe and classify) and the LMC (where the main selection effects are the brightness of the X-ray source and our ability to actually identify the counterpart in crowded fields). Therefore we conclude that the observed spectral distribution is not dominated by selection effects. The only selection effect that could affect both samples would be the existence of a population of objects with very low X-ray luminosities, but such a population should be evident in the Solar neighbourhood, if it was a major contributor to the numbers of Be/X-ray binaries.  \\subsection{Implications for binary evolution}  Massive X-ray binaries are born as the result of the evolution of a massive binary in which mass transfer between the two components has taken place before one of them (originally the most massive) became a compact object after gravitational collapse of the core. In what follows we will adopt the following convention: in the massive X-ray binary, we will refer as the {\\em donor} to the massive star which is passing mass to the {\\em compact object} (which will always be referred to as such); in the original binary before the gravitational collapse, we will call the {\\em original primary} to the component which evolved to become the compact object (and which was the more massive component when the system reached the ZAMS) and the {\\em original secondary} to the component that would later become the donor.   \\begin{table}[ht] \\caption{Other suggested X-ray binaries in the LMC, not included in this sample. References are Cowley et al. (1997) = C97, Haberl \\& Pietsch (1999) = HP99 and Sasaki et al. (2000) = SHP00.} \\begin{center} \\begin{tabular}{lcc} \\hline Name & Suggested & Reference\\\\ & Counterpart & \\\\ \\hline \\object{RX J0516.0$-$6916} & B1V & C97 \\\\ \\object{RX J0512.6$-$6717} & $-$ & HP99\\\\ \\object{RX J0532.4$-$6535} & {\\object GRV\\,0532$-$6536} & HP99\\\\ \\object{RX J0535.8$-$6530} & $-$ & HP99\\\\ \\object{RX J0541.4$-$6936} & \\object{Sk $-69\\degr$ 271} & SHP00\\\\ \\object{RX J0541.6$-$6832} & \\object{BI 267}& SHP00\\\\ \\hline \\end{tabular} \\end{center} \\label{tab:unknowns} \\end{table}   Be/X-ray binaries are generally believed to be formed via a standard evolutionary channel. The progenitor is an intermediate-mass close binary with moderate mass ratio $q\\ga 0.5$. The original primary starts transfering mass to its companion after the end of the hydrogen core burning phase (case B), resulting in a helium star and a rejuvenated main sequence star. If the helium star is massive enough, it will undergo a supernova explosion and become a neutron star. If the binary is not disrupted, it can then become a Be/X-ray binary. This model has been developed by Habets (1987), Pols et al. (1991), Portegies Zwart (1995) and Van Bever \\& Vanbeveren (1997), who have calculated the expected population distribution when different assumptions are made. All these models assume that the original primary must have a mass $\\ga 12\\, M_{\\sun}$ (because otherwise it would not produce a neutron star) and that strong tidal interactions during the mass transfer phase result in the circularization of the orbit. It is implicitly assumed that the Be nature of the original secondary is due to accretion of high-angular-momentum material from the primary, even though no description of the exact physical process has been attempted.  Habets (1987) studied possible evolutionary scenarios and came to the conclusion that, if supernova explosions are always symmetric, only low-eccentricity Be/X-ray binaries may exist, because the exploding helium star has in all cases a much lower mass than its companion and therefore only a small fraction of the system mass is lost in the explosion. As a consequence, van den Heuvel \\& van Paradijs (1997) conclude that the observational detection of Be/X-ray binaries with very eccentric orbits is proof of the existence of intrinsic kicks imparted to the neutron stars during the collapse that leads to their formation.  A second evolutionary channel has been explored by Habets (1987). In this case, the Be/X-ray binary is formed as the result of the evolution of a binary containing a massive star ($M_{{\\rm MS}} \\ga 20\\,M_{\\sun}$ where $M_{{\\rm MS}}$ is the main sequence mass of the primary) and a rather less massive companion (of mass $M_{*} \\sim 10\\,M_{\\sun}$). Because of the $q \\la 0.5$ mass ratio, when the primary fills its Roche lobe, mass transfer is highly non-conservative, a common envelope forms and essentially all of the hydrogen-rich envelope of the primary is lost from the system. The resulting system consists of a relatively massive He star ($M_{{\\rm He}} \\sim 5\\,-\\,10\\: M_{\\sun}$) and an unaffected secondary. Since the helium star still contains a substantial fraction of the total binary mass, a symmetric supernova explosion can now result in a neutron star orbiting a Be star in a very eccentric orbit, ($e = M_{{\\rm lost}}/M_{{\\rm left}}$, where $M_{{\\rm lost}}$ is the mass lost from the system during the explosion and $M_{{\\rm left}}= M_{*} + M_{{\\rm x}}$ is the mass of the two components after the explosion).  An important limitation found by Habets (1987) is that, because of the initial conditions required, this channel may only explain Be/X-ray binaries in which a neutron star orbits a low-mass Be star in a close, eccentric orbit. This channel cannot produce systems with either massive donors ($M_{*} \\ga 15\\, M_{\\sun}$) or wide orbits. Also important is the fact that, since the original secondary is left basically unaffected by the whole process, it has to be implicitly accepted that it becomes a Be star for some {\\em intrinsic} reason, unlike in the first channel considered.  The spectral distribution found for Be/X-ray binaries in both the Milky Way and the LMC, sharply peaked at spectral type B0 (roughly corresponding to $M_{*}\\approx 16\\,M_{\\sun}$) and not extending beyond B2 ($M_{*}\\approx 10\\,M_{\\sun}$) strongly argues against the second channel contributing significantly to the formation of Be/X-ray binaries. It is therefore strong indirect evidence in support of the existence of supernova kicks.  In spite of this, the second channel is not only physically possible, but almost certainly realized: a system like LMC X-3 (in which the donor is {\\em not} a Be star) must have formed through such a process, since the original primary must have been sufficiently massive to result in a black hole and the original secondary cannot have been more massive than the present $M_{*}\\la 8\\,M_{\\sun}$. Therefore, in the original system, $q\\la0.3$ and non-conservative evolution must have occurred -- perhaps through case C mass transfer, as in the models by Brown et al. (2001), since Wellstein \\& Langer (1999) argue that cases A and B cannot result in black holes.  Though it is not impossible that some of the Be/X-ray binaries with lower-mass donors may have formed through the second channel, it is clear that most must have followed the first channel. Taken at face value, the dearth (or even absence) of Be/X-ray binaries forming through the second channel could be interpreted as a suggestion that a star does not develop Be characteristics due to intrinsic reasons, but must have undergone a process of mass transfer in a binary, as discussed by Gies (2000). Such scenario is in disagreement with the population synthesis models of Portegies Zwart (1995) and van Bever \\& Vanbeveren (1997), which indicate that not all Be stars may have formed through binary evolution and therefore support an intrinsic cause for the Be phenomenon. The dominance of the first channel may then simply be a reflection of a preference for binary systems which originally have a mass ratio $q\\approx1$. If the progenitor binaries that will evolve through the second channel are much less numerous than those which follow the first, their end products will naturally be also a minority.  It is, however, intriguing to note that, while {\\em no} system containing a Be star and a black hole is known, several systems containing (non-emission) OB stars and a black hole are known, and their donors span a wide range of spectral types. Since their number is still relatively low, one can attribute the absence of Be + black hole binaries to low number statistics. This is, in principle, a valid argument, but it must be noted that black hole binaries with a Be companion should have a much longer lifetime as an X-ray source than any OB + black hole binary, because a mechanism for mass transfer (the Be disk) exists during a longer period. The situation should not be very different from binaries containing a neutron star, where close to 70\\% of systems known are Be/X-ray binaries. Therefore the absence of known Be + black hole binaries strongly suggests that there is some physical reason why Be + black hole binaries cannot be formed or are not bright X-ray sources."
    },
    "0201/astro-ph0201498_arXiv.txt": {
        "abstract": "We introduce a new cosmological diagnostic pair $\\lbrace r,s\\rbrace$ called Statefinder. The Statefinder is a geometrical diagnostic and allows us to characterize the properties of dark energy in a model independent manner. The Statefinder is dimensionless and is constructed from the scale factor of the Universe and its time derivatives only. The parameter $r$ forms the next step in the hierarchy of geometrical cosmological parameters after the Hubble parameter $H$ and the deceleration parameter $q$, while $s$ is a linear combination of $q$ and $r$ chosen in such a way that it does not depend upon the dark energy density. The Statefinder pair $\\lbrace r,s\\rbrace$ is algebraically related to the equation of state of dark energy and its first time derivative. The Statefinder pair is calculated for a number of existing models of dark energy having both constant and variable $w$. For the case of a cosmological constant the Statefinder acquires a particularly simple form. We demonstrate that the Statefinder diagnostic can effectively differentiate between different forms of dark energy. We also show that the mean Statefinder pair can be determined to very high accuracy from a SNAP-type experiment. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0201/hep-th0201066_arXiv.txt": {
        "abstract": "We apply methods of dynamical systems to study the behaviour of the Randall-Sundrum models. We determine evolutionary paths for all possible initial conditions in a 2-dimensional phase space and we investigate the set of accelerated models. The simplicity of our formulation in comparison to some earlier studies is expressed in the following: our dynamical system is a 2-dimensional Hamiltonian system, and what is more advantageous, it is free from the degeneracy of critical points so that the system is structurally stable. The phase plane analysis of Randall-Sundrum models with isotropic Friedmann geometry clearly shows that qualitatively we deal with the same types of evolution as in general relativity, although quantitatively there are important differences. ",
        "introduction": "The rapid development of particle-physics-motivated cosmology represented mainly by superstring cosmology \\cite{ven91,gv93,superjim} has led to a change of views onto the standard problems of inflationary cosmology such as the (past) horizon and flatness problems \\cite{guth,linde}. On the other hand, the astronomical observations of supernovae Ia \\cite{Supernovae} strongly suggest that the universe not only had possibly accelerated at its early stages of evolution but it is also accelerating now. This puts strong constraints onto the matter content of the universe because only the exotic (negative pressure) matter in standard cosmology can lead to acceleration. However, a phenomenological nature of such an exotic matter in standard models expressed in terms of the perfect fluid does not seem to make enough connection with particle physics and that is why one is looking for other, more physical, descriptions. Among them the main proposal is quintessence or time-dependent scalar field \\cite{steinhardt} which substitutes an ordinary phenomenological barotropic fluid. However, there are other interesting proposals to express the phenomenon of acceleration which are related to superstring or M-theory models. The best known are pre-big-bang \\cite{ven91,gv93}, brane \\cite{hw,rs1,rs2} and ekpyrotic models \\cite{turok1,turok2}. In this paper we try to study the standard cosmology problems (such as cosmic acceleration and past horizon problems) within the framework of the brane universes. We do not study, for instance, the future horizon problem following the recent discussion inspired by the S-matrix formulation within the superstring theories \\cite{futurehor}. In the paper we are interested mainly in the early stage of the evolution of the universe although the supernovae data gives restrictions onto this stage and they should be taken into account.  The idea of brane universes has been first presented by Ho\\v{r}ava and Witten \\cite{hw} who considered strong coupling limit of heterotic $E_8 \\times E_8$ superstring theory, i.e., M-theory. This limit results in `exotic' \\cite{Visser85,Barcelo00} Kaluza-Klein type compactification of $N = 1$, $D =11$ supergravity on a $S^1/Z_2$ orbifold (a unit interval) in a similar way as compactification of $N=1$, $D=11$ supergravity on a circle $S^1$ results in strongly coupled limit of type IIA superstring theory. In Ho\\v{r}ava-Witten theory there exist two 10-dimensional branes to which all the gauge interactions are confined, and they are connected via the orbifold, with gravity propagating in all 11 dimensions. After further compactification of Ho\\v{r}ava-Witten models on a Calabi-Yau manifold one gets an effective 5-dimensional theory which has been applied to cosmology \\cite{lukas,reall,jim,mpd00}.  Randall and Sundrum \\cite{rs1,rs2} developed similar to Ho\\v{r}ava-Witten scenario which was mainly motivated by the hierarchy problem in particle physics \\cite{Arkani98,Arkani99}. As a result, they obtained a 5-dimensional spacetime (bulk) with $Z_2$ symmetry with two/one 3-brane(s) embedded in it to which all the gauge interactions are confined. In one-brane scenario \\cite{rs2} the brane appears at the $y=0$ position, where $y$ is an extra dimension coordinate and the 5-dimensional spacetime is an anti-de Sitter space with negative 5-dimensional cosmological constant. The extra dimension can be infinite due to the exponential `warp' factor in the metric \\begin{equation} \\label{rsmet} ds^2 = \\exp{\\left(-2\\frac{\\mid y \\mid}{l}\\right)} \\left[ -dt^2 + d\\vec{x}^2 \\right] + dy^2, \\end{equation} and $l$ gives the curvature scale of the anti-de Sitter space. In the simplest case the induced metric on a brane is a Minkowski metric (energy momentum tensor of matter vanishes). However, the requirement to allow matter energy-momentum tensor on the brane leads to breaking of conformal flatness in the bulk, and the metric (\\ref{rsmet}) is no longer valid. This fact is obviously related to the appearance of the Weyl curvature in the bulk \\cite{BDL,BDEL,roy01b}. The full set of 5-dimensional and projected 4-dimensional equations has been presented in Refs.~\\cite{Shiromizu00,Sasaki00,Mukhoyama00}. Global geometric properties of such brane models have also been studied (see Ref. \\cite{Ishihara01}). Generalized bulk spacetimes different from those of (\\ref{rsmet}) have been found in Ref. \\cite{roy01a}. Anisotropic Kasner branes have been immersed in AdS bulk in Ref.~\\cite{afrolov}.  Campos and Sopuerta \\cite{Campos01a,Campos01b} used the dynamical system methods to the analysis of the Friedmann, Bianchi I and V Randall-Sundrum brane world type cosmological models with a non-vanishing 5-dimensional Weyl tensor. They considered the dynamics of this model in the form of a {\\it higher-than-two-dimensional} dynamical system. Exact analytic brane configurations with a vanishing Weyl tensor with perfect and viscous fluid have been presented in Refs. \\cite{harko1,harko2}. Coley \\cite{Coley01a,Coley01b} also studied the dynamics of these Randall-Sundrum models - he made a step towards Mixmaster (Bianchi IX) dynamics and found it was not chaotic provided the matter had positive pressure. In this paper we show that the dynamical system which describes the evolution of the brane models (both isotropic (FRW) and anisotropic (Bianchi I or Bianchi V) types) can be represented in the simplest way in the form of a {\\it two-dimensional\\/} Hamiltonian dynamical system. Such visualization has a great advantage because we {\\it avoid\\/} the problem of degeneracy of critical points which appeared in the {\\it higher dimensional} phase space (see Ref.~[13] in \\cite{Campos01b}). It is well known that the existence of such critical points is a possible reason for the structural instability of a model. On the other hand, the representation of dynamics as a one-dimensional Hamiltonian flow allows to make the classification of possible evolution paths in the {\\it configuration space} which is complementary to phase diagrams. It also makes simpler to discuss the physical content of the model. Finally, the construction of the Hamiltonian allows to study quantum cosmology on the brane as it was attempted in Refs. \\cite{nunez1,nunez2}, in full analogy to what is usually done in general relativity \\cite{DL95}.  In this paper we demonstrate the effectiveness of representation of dynamics as a one-dimensional Hamiltonian flow. In this representation the phase diagrams in a two-dimensional phase space allow to analyze the {\\it acceleration and (past) horizon} problems in a clear way. We reduce the dynamics to a two-dimensional phase space with an autonomous system of equations.  We deal with the full global dynamics of brane models whose asymptotic states are represented by critical points of the system. From theoretical point of view it is important how large the class of accelerated models is. We will call this class of accelerated models typical/generic, if the domain of acceleration in the phase space is a non-zero measure. On the other hand, if only the non-generic (zero-measure) trajectories are represented by accelerated models, then the mechanism which drives these trajectories should be called ineffective. Such a point of view is a consequence of the fact that, if the acceleration is an attribute of a trajectory which starts with a given initial conditions, it should also be an attribute of a trajectory which starts with a neighbouring initial conditions.  Our analysis of the brane-world type model may be considered as complementary to Campos and Sopuerta's analysis \\cite{Campos01a,Campos01b} (see also \\cite{Coley01a,Coley01b}). However, it is more advantageous in many points. It is because our dynamical system is a 2-dimensional system which is the smallest possible dimension to study isotropic cosmological systems of equations. This allows to avoid huge redundancy of degenerated critical points and trajectories which appear in a higher dimensional phase space (cf. phase diagrams of \\cite{Campos01a,Campos01b}) so to the structural instability. Finally, for a two-dimensional Hamiltonian system one is able to study all its properties in a configuration space rather than in a phase space.  The paper is organized as follows. In Section 2 we present simple Hamiltonian dynamics of the brane universes. In Section 3 we discuss cosmological models in configuration space, stability and the domains of the phase plane for which cosmic acceleration and horizon problem avoidance appear. In Section 4 we present phase portraits for FRW models with dust ($\\gamma =1$) and domain-wall-like matter ($\\gamma = 1/3$) \\cite{Supernovae,stelbd93,Dabrowski96} which induces accelerated expansion (quintessence) in standard general relativistic cosmology. In section 5 we present simple form of dynamical systems in which constant coefficients play the role of the observational dimensionless density of matter $\\Omega$ parameters. In Section 6 we discuss our results. ",
        "conclusions": "In the paper we studied the dynamics of Randall-Sundrum brane universes with isotropic Friedmann geometry. Our approach is the simplest in that we formulate the problem in a 2-dimensional phase space and not in a higher-dimensional phase space as it has been done recently. First of all, this is the smallest possible dimension to study the isotropic cosmological systems of equations. Then, it allows to {\\it avoid} the degeneracy of critical points and trajectories which appear in a higher dimensional phase space and so to {\\it avoid} structural instability of these models. Also, since the system is Hamiltonian one is able to study its properties in a one-dimensional configuration space rather than in phase space. In such a space the motion of the universe point is reduced to the motion of a particle with the potential energy $V(a)$, where $a$ is the scale factor (see Section 2). On the other hand, Hamiltonian formulation can easily be applied to quantum cosmology on the brane.  In Fig.~\\ref{fig:1} and \\ref{fig:2} we can see that there is no qualitative difference in the shapes of the potential function $V(a)$ as $\\gamma$ varies for cases under studies. Only the values of $a$ for which the potential is negative have the physical meaning. The maximum of the potential function corresponds to an unstable saddle point on a phase diagram (a critical point solution of the zero energy level). The de Sitter solution is a point at infinity. Near the initial singularity we have a solution $a(t) \\propto t^{1/3\\gamma}$ which differs from a general relativistic one $a(t) \\propto t^{2/3\\gamma}$ .  From the theory of qualitative differential equations in Sections 3 and 4 we obtained the visualization of the system evolution in the phase plane $(x,\\dot{x})$ and analyzed the asymptotic states and concluded that the brane models are {\\it structurally stable}. We studied the solution of the (past) horizon problem and the initial conditions for acceleration of our Universe. We have the neat interpretation of {\\it a domain of acceleration} as a domain in configuration space where the potential function decreases. Therefore from the observation of the potential function $V(a)$, we can see the acceleration domain $a>a_{\\text{max}}$, with $V'(a_{\\text{max}})=0$, which is independent of the curvature index $k$. On the other hand, the (past) horizon problem is solved when $V(a) \\to \\text{const}$ (may be equal to zero) as $a \\to 0$. If we find trajectories for which $y = \\dot{a}$ goes to infinity as $x = a \\to 0$, then the horizon is present in such a model. This is the case of dust brane model as well as a brane world filled with domain-wall-like matter with $\\gamma = 1/3$.  In general our model does not solve the horizon problem due to the existence of matter dominated phase of the early evolution where the term $\\rho^2$ dominates and so $H \\propto \\rho \\propto a^{-3\\gamma}$ ($a(t) \\propto t^{-1/3\\gamma}$). Only in the special case of heavy domain-wall-like matter ($\\gamma \\le 1/3$) the horizon problem is solved, because the evolution near the singularity dominated by matter (brane tension) is $\\dot{a} \\propto a^{1-3\\gamma}$. However, the existence of such a singularity requires ${\\cal U} =0$ (Weyl tensor contribution from the bulk vanishes). Therefore, the case $\\gamma \\le 1/3$ is strictly distinguished because as $x \\to 0$ then $y \\to \\text{const}$. For $\\gamma = 1/3$ the term $\\rho^{2}/\\lambda$ vanishes and brane effects are negligible.  In the phase portraits we can observe similarities as well as differences in both considered cases. In the dust-like matter case there is a quasi-static stage first discussed by Lema{\\^\\i}tre (called `loitering stage' in \\cite{Sahni92}). Such quasi-static stages, present in the case of wall-like-matter, corresponds to configurations in the vicinity of a critical point. The acceleration does not depend on the value of $\\Omega_{k}$, but in the phase space there are different sets of initial conditions which provide acceleration.  We conclude that the classes of admissible models in brane scenario are {\\it qualitatively} the same as in general relativity, i.e., there exists a homeomorphism in the phase space which transforms trajectories of brane models into general relativistic ones without a change of orientation. However, {\\it quantitatively} the models differ from general relativistic ones - for instance, they asymptote the singularity in a different way.  Finally, we reformulated the phase space quantities interms of the observational parameters $\\Omega$ and presented the formalism in which one is able to perform the observational tests (redshift-magnitude relation) of brane models which we will address in a separate paper."
    },
    "0201/astro-ph0201337_arXiv.txt": {
        "abstract": "\\noindent {\\it We show that the astrometric Hipparcos data of the stars hosting planet candidates are not accurate enough to yield statistically significant orbits. Therefore, the recent suggestion, based on the analysis of the Hipparcos data, that the orbits of the sample of planet candidates are not randomly oriented in space, is not supported by the data. Assuming random orientation, we derive the mass distribution of the planet candidates and shows that it is flat in log M, up to about 10 \\MJ. Furthermore, the mass distribution of the planet candidates is well separated from the mass distribution of the low-mass companions by the 'brown-dwarf desert'. This indicates that we have here two distinct populations, one which we identify as the giant planets and the other as stellar secondaries.  We compare the period and eccentricity distributions of the two populations and find them surprisingly similar.  The period distributions between 10 and 1650 days are flat in log period, indicating a scale-free formation mechanism in both populations. We further show that the eccentricity distributions are similar --- both have a density distribution peak at about 0.2--0.4, with some small differences on both ends of the eccentricity range. We present a toy model to mimic both distributions. The toy model is composed of Gaussian radial and tangential velocity scatters added to a sample of circular Keplerian companions. A scatter of a dissipative nature can mimic the distribution of the eccentricity of the planets, while scatter of a more chaotic nature could mimic the secondary eccentricity distribution. We found a significant paucity of massive giant planets with short orbital periods.  The low-frequency of planets is noticeable for masses larger than about 1 \\MJ\\ and periods shorter than 30 days. We point out how, in principle, one can account for this paucity. } ",
        "introduction": "More than fifty candidates for extrasolar planets have been announced over the past six years (e.g., Schneider 2001).  In each case, precise stellar radial-velocity measurements indicated the presence of a low-mass unseen companion, with a minimum mass between 1 and about 10 Jupiter masses (\\MJ). The identification of these unseen companions as planets relied on their masses being in the planetary range.  However, the actual masses of the planet candidates are not known. The radial-velocity data yield only $M_2 \\sin\\, i$, where $M_2$ is the secondary mass and $i$ is the inclination angle of its orbital plane, which cannot be derived from the spectroscopic data. Nevertheless, the astronomical community considered the planet-candidate masses as being close to their derived minimum masses --- $M_2 \\sin\\, i$. This is so because at random orientation the most probable inclination is $90\\degrm$, and the expected value of $\\sin\\,i$ is close to unity.  Very recently some doubt has been cast about the validity of the random orientation assumption. \\GHB\\ and \\HBG\\ analysed the Hipparcos astrometric data of the stars hosting planet candidates {\\it together} with the stellar precise radial-velocity measurements and derived in some cases very low inclination angles for the orbital planes. \\HBG\\ found eight out of 30 systems with an inclination smaller or equal to $0.5^\\circ$, four of which they categorized as highly significant. The probability of finding such small inclinations in a sample of orbits that are {\\it isotropically} oriented in space is extremely small, indicating either a problematic derivation of the astrometric orbit, or, as suggested by \\HBG, some serious orientation bias in the inclination distribution of the sample of detected planet candidates.  However, the analysis of the Hipparcos data can be misleading.  As has been shown by Halbwachs et al.\\ (2000), one can derive a small {\\it false} orbit with the size of the typical positional error of Hipparcos, about 1 milli-arc-second (=\\mas), caused by the scatter of the individual measurements. Therefore, one should carefully evaluate the statistical significance of any astrometric orbit of that size derived from the Hipparcos data. In Section~2 we summarize our work (Zucker \\& Mazeh 2001a) that evaluates the significance of the astrometric orbits by applying a permutation test to the Hipparcos data. Similarly to the results of Pourbaix (2001) and Pourbaix \\& Arenou (2001), we also find that the significance of all the Hipparcos astrometric orbits of the planet candidates are less than 99\\%, including \\RCB\\ that attracted much attention after the publication of \\GHB\\ suggestion. We therefore conclude that the Hipparcos data does not prove the anisotropy of the orientations of the orbital planes of the planet candidates.  After showing that the random orientation in space is still a reasonable assumption, not confronted by any available measurement, we present in Section 3 our work (Zucker \\& Mazeh 2001b) that uses this assumption to derive the mass distribution of the planet candidates. This is done with \\MAXLIMA, a MAXimum LIkelihood MAss algorithm which we constructed to derive the mass distribution. Similar to the results of Jorrisen, Mayor \\& Udry (2001), we show that the mass distribution of the planet candidates is separated from the one of the secondary masses by the so-called 'brown-dwarf desert' (e.g., Marcy \\& Butler 2000). This indicates that we are dealing with two different classes of objects. One is the giant planets, with masses not far from the planetary mass range, while the other is the low-mass secondaries, with stellar mass range.  One could speculate that the separation between the two different mass distributions indicates different formation processes.  The commonly accepted paradigm is that planets were probably formed by coagulation of smaller, possibly rocky, bodies, whereas stars were probably formed by some kind of fragmentation of larger bodies. In other words, planets were formed by small bodies that grew larger, whereas stars, binary included, were formed by fragmentation of large bodies into smaller objects (e.g., Lissauer 1993; Black 1995).  This could imply, for example, that the distribution of orbital eccentricities of giant planets and low-mass binaries would be substantially different.  All the solar planets have nearly circular orbits, whereas binaries have eccentric orbits (e.g., Mazeh, Mayor, \\& Latham 1996). We could also expect the periods of planets to be longer than 10 years, like the giant planets in the solar system. Many studies of the newly discovered planets showed that this is not the case (e.g., Marcy, Cochran, \\& Mayor 2000). Moreover, following Heacox (1999) who based his analysis upon only 15 binaries and a handful of planet candidates, we show in Section 4 that within some reasonable restrictions, the eccentricity and period distributions of the two samples are surprisingly similar.  Similar results have been obtained by Stepinski \\& Black (2001a,b,c).  In Section 5 we consider a toy model that can generate the eccentricity distribution of both populations. ",
        "conclusions": "The logarithmic mass distribution derived here shows that the planet candidates are indeed a separate population, probably formed in a different way than the secondaries in spectroscopic binaries. Surprisingly the eccentricity and period distributions, with some restrictions, are very much the same.  Furthermore, the two period distributions follow strictly a straight line. This indicates flat density distributions on a logarithmic scale, inconsistent with the Duquennoy \\& Mayor (1991) log-normal distribution. Interestingly, flat logarithmic distribution is the only scale-free distribution, and could be argued to be the most simple distribution. Maybe the two populations were formed by two different mechanisms that still have this scale-free feature in common (Heacox 1999).  The eccentricity distribution of the sample of giant planets and that of stellar companions are similar (Stepinski \\& Black 2001c).  In spite of the similarity, they are not identical, especially if compared to the remarkable similarity between the two period distributions.  The eccentricity distributions can be attained by Keplerian orbits whose velocities are normally disturbed in the tangential and the radial directions.  We found a significant paucity of large planets with short orbital periods, and point out how, in principle, one can account for this paucity."
    },
    "0201/astro-ph0201101_arXiv.txt": {
        "abstract": "{ The metal line profiles of different ions observed in high \\ion{H}{i} column density systems [$N$(\\ion{H}{i}) $> 10^{16}$ cm$^{-2}$] in quasar spectra can be used to constrain the ionization structure and kinematic characteristics of the absorbers. For these purposes, a modified Monte Carlo Inversion (MCI) procedure was applied to the study of three absorption systems in the spectrum of the HDF-South quasar J2233--606 obtained with the UVES spectrograph at the VLT/Kueyen telescope. The MCI does not confirm variations of metal abundances within  separate systems which were discussed in the literature. Instead, we found that an assumption of a homogeneous metal content and a unique photoionizing background is sufficient to describe the observed complex metal profiles. It was also found that the linear size $L$ and the line-of-sight velocity dispersion $\\sigma_{\\rm v}$ measured within the absorbers obey a scaling relation, namely, $\\sigma_{\\rm v}$ increases with increasing $L$, and that metal abundance is inversely proportional to the linear size of the system: the highest metallicity was measured in the system with the smallest $L$. ",
        "introduction": "\\addtocounter{footnote}{3}  Absorption systems in quasar spectra provide unique information on the intervening intergalactic matter (IGM) up to redshift $z \\simeq 6$, back to the time when the Universe was less than 7\\% of its present age. High resolution spectroscopic observations available nowadays at large telescopes open new opportunities to investigate the physical nature of quasar absorbers. Reliable data on the chemical composition of the IGM and on the physical characteristics (like velocity and density distributions, volumetric gas density, kinetic temperature, ionization structure etc.)  of the absorbers is an important clue to our understanding of galaxy formation, chemical evolution of the IGM and the origin of the large-scale structure.  In recent investigations much attention find the metal systems  which are the absorbers exhibiting as a rule numerous lines of low (like \\ion{H}{i}, \\ion{C}{ii}, \\ion{Si}{ii},  \\ion{Mg}{ii}, \\ion{Al}{ii}) and high (like \\ion{C}{iii}, \\ion{N}{iii}, \\ion{Si}{iii}, \\ion{C}{iv}, \\ion{Si}{iv}, \\ion{N}{v}) ionized species.  Presence of metals provides a unique opportunity to study the physical conditions of matter at early epochs. Unfortunately, the computational methods usually applied to high resolution spectra lie quite often  behind the quality of observational data and fail to extract from them all encoded information. The common processing method consists of the deconvolution of complex absorption profiles into an arbitrary number of separate components (assuming a constant gas density within each of them) which are then fitted to Voigt profiles. However, in many cases this procedure may not correspond to real physical conditions: observed complexity and non-similarity of the profile shapes of different ions indicate that these systems are in general absorbers with highly fluctuating density and velocity fields tightly correlated with each other. Too high or too low gas temperatures, extremely varying metallicities between subcomponents, exotic UV background spectra and other physically badly founded outcomings may be artifacts arising from the inconsistency of the applied methods (see examples in Levshakov et al. 1999; Levshakov et al. 2000b, hereafter Paper~I).  In Paper~I we developed a new method for the QSO spectra inversion, -- the Monte Carlo Inversion (MCI), -- assuming that the absorbing region is a cloud with uniform metallicity but with fluctuating density and velocity fields inside it. This computational procedure which is based on stochastic optimization allows us to recover both the underlying hydrodynamical fields and the physical parameters of the gas. First application of the MCI to the analysis of the $z_{\\rm abs} = 3.514$ system toward Q08279+5255 (Levshakov et al. 2000a) has shown that the proposed method is very promising especially in the inversion of complex absorption spectra with many metal lines.  In this paper we start a new comprehensive survey of the metal systems for which high resolution and high signal-to-noise spectra are available. We present here the results for three absorption systems ($z_{\\rm abs}$ = 1.87, 1.92 and 1.94) from the spectrum of the quasar J2233--606 which have been already studied by Prochaska \\& Burles (1999), and D'Odorico \\& Petitjean (2001, hereafter DP) using the common Voigt fitting method. We re-calculate these systems using the MCI in order to compare the applicability of both approaches and to show up their restrictions.  The structure of the paper is as follows. In Sect.~2 the data sets used in the MCI analysis are described. Sect.~3 contains the details of the applied computational procedure. The results obtained for each of the mentioned above systems are presented in Sect.~4. Conclusions are reported in Sect.~5. In Appendix the general equation of the entropy production rate is given which is used to calculate the density and velocity configurations along the line of sight exhibiting minimum dissipation.   \\begin{figure*} \\vspace{0.0cm} \\hspace{0.0cm}\\psfig{figure=H3203F1.PS,height=13.0cm,width=20.0cm} \\vspace{-0.5cm} \\caption[]{ Hydrogen and metal absorption lines associated with the $z_{\\rm abs} \\simeq 1.92$ system toward J2233--606 (normalized intensities are shown by dots with $1\\sigma$ error bars). The zero radial velocity is fixed at $z = 1.92595$. Smooth lines are the synthetic spectra convolved with the corresponding point-spread functions (FWHM$_{\\rm VLT} = 6.7$ km~s$^{-1}$, FWHM$_{\\rm HST} = 10.0$ km~s$^{-1}$) and computed with the physical parameters from Table~1. Bold horizontal lines mark pixels included in the optimization procedure. The normalized $\\chi^2_{\\rm min} = 1.10$ (the number of degrees of freedom $\\nu = 1450$) } \\label{fig1} \\end{figure*} ",
        "conclusions": "A main goal of this work was to investigate the reliability of the physical parameters and dynamical characteristics of the metal absorption systems obtained by means of the standard Voigt fitting procedure and by the modified Monte Carlo Inversion algorithm. For comparison we used the recently obtained results on the Voigt fitting analysis of three systems toward J2233--606 (D'Odorico \\& Petitjean 2001).  We found that both approaches deliver similar total column densities of unsaturated metal lines. The saturated profiles may, however, be treated differently (e.g., the \\ion{Si}{iii} $\\lambda1206$ \\AA\\, line in the $z = 1.94$ system).  We also found that metal abundances based on the Voigt deconvolution procedure differ considerably from those obtained by the MCI. For instance, instead of fluctuating metallicities found in the absorbers toward J2233--606 by DP, the MCI shows that an assumption of a homogeneous metal abundance for the whole system under study is quite sufficient to represent all observed features.  New and principal results which can be obtained only with the MCI procedure are the kinematic characteristics of the absorbers. We estimated selfconsistently for the first time the density and velocity dispersions along the sightlines within the absorbers and calculated the total hydrogen column and volumetric densities which gave us a direct measure of their linear sizes. The found dimensions of $L \\simeq 5$ kpc to $L \\simeq 80$ kpc are in good agreement with measurements of the extended gaseous envelopes around the nearby galaxes which were probed by the \\ion{Mg}{ii} absorption lines (Bergeron \\& Boiss\\'e 1991) and by the \\ion{C}{iv} lines (Chen et al. 2001).  A new issue obtained in the MCI analysis is a scaling relation. Namely, we found that the linear size $L$ shows a positive correlation with the line-of-sight velocity dispersion $\\sigma_{\\rm v}$, i.e. the higher $L$, the larger $\\sigma_{\\rm v}$ is observed (see Table~1). Although our sample is still too small to carry out statistical analysis of this correlation, the scaling tendency is of the same kind that can be expected for virialized systems. The velocity dispersion is closely related to the total mass of the system in a stationary state (cf. the scaling law known as the {\\it fundamental plane} for elliptical galaxies). Taking into account that the scaling laws are different for different types of objects (see, e.g., Fig.~2 in Mall\\'en-Ornelas et al. 1999), future statistical analysis may allow us to classify absorbers at different redshifts.  It is also interesting to note another scaling relation: we observe systematically higher metal abundance with decreasing $L$, and vice verse, the higher $L$, the lower metallicity is deduced. If $L$ reflects the linear size of a distant absorber, then we may conclude that a compact absorber has, presumably, higher metal content as compared with an extended one."
    },
    "0201/astro-ph0201271_arXiv.txt": {
        "abstract": "We present a numerical investigation of dead, or relic, radio galaxies and the environmental impact that radio galaxy activity has on the host galaxy or galaxy cluster.  We perform axisymmetric hydrodynamical calculations of light, supersonic, back-to-back jets propagating in a $\\beta$-model galaxy/cluster atmosphere.  We then shut down the jet activity and let the resulting structure evolve passively.  The dead source undergoes an initial phase of pressure driven expansion until it achieves pressure equilibrium with its surroundings.  Thereafter, buoyancy forces drive the evolution and lead to the formation of two oppositely directed plumes that float high into the galaxy/cluster atmosphere.  These plumes entrain a significant amount of low entropy material from the galaxy/cluster core and lift it high into the atmosphere.  An important result is that a large fraction (at least half) of the energy injected by the jet activity is thermalized in the ISM/ICM core.  The whole ISM/ICM atmosphere inflates in order to regain hydrostatic equilibrium.  This inflation is mediated by an approximately spherical disturbance which propagates into the atmosphere at the sound speed.  The fact that such a large fraction of the injected energy is thermalized suggests that radio galaxies may have an important role in the overall energy budget of rich ISM/ICM atmospheres.  In particular, they may act as a strong and highly time-dependent source of negative feedback for galaxy/cluster cooling flows. ",
        "introduction": "There is now overwhelming observational and theoretical evidence that powerful radio galaxies possess highly collimated and relativistic twin jets of matter that emerge from a central active galactic nucleus (AGN).  In the powerful Faranoff-Riley type II radio galaxies (FR-II; Fanaroff \\& Riley 1974), the jets are thought to remain relativistic along most of their length before terminating in a series of shocks resulting from interaction with the surrounding material.  A substantial amount of theoretical and numerical (hydrodynamic) work suggests that, after passing through the terminal shock, the spent jet material inflates a broad cocoon which encases the jets (e.g., see Scheuer 1974; Norman et al. 1982; Lind et al. 1989; Begelman \\& Cioffi 1989; Cioffi \\& Blondin 1992; Hooda, Mangalam \\& Wiita 1994).  In the early life of a powerful source, 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.  From the onset of activity, Kelvin-Helmholtz (KH) instabilities act to shred the contact discontinuity.  If the ambient medium has a density profile characterized by $\\rho\\propto r^{-\\alpha}$ where $\\alpha>2$, the contact discontinuity will accelerate and hence will be Rayleigh-Taylor (RT) unstable.  However, in the more physical case of $\\alpha<2$ the contact discontinuity will possess an initial deceleration, which is sufficient to stabilize the contact discontinuity against RT modes.  Only at later times, once the cocoon comes into approximate pressure balance with its surroundings, will the deceleration of the contact discontinuity no longer exceed the local gravitational acceleration, thereby rendering it RT unstable.  Eventually, AGN activity will cease and the jets will turn off.  Although this stage of a radio galaxy's life has been little studied, it seems likely that the combined action of the KH and RT instabilities will transform the remnant of the cocoon into `plumes' that rise in the potential of the galaxy/cluster under the action of buoyancy forces (Gull \\& Northover 1973; Churazov et al. 2000; Br\\\"uggen \\& Kaiser 2000).  Observationally, radio galaxies are often seen to be associated with galaxy or cluster cooling flows.  This raises an obvious and interesting set of questions.  To what extent is radio galaxy activity a natural {\\it response} to the cooling flow phenomenon?  Do radio galaxies act back on their environment to partially offset the cooling flow, i.e., are radio galaxies a dramatic manifestation of le Chatelier's principle?  Is radio galaxy activity a crucial component for our understanding of galaxy clusters?  Since it takes only a small fraction of the cooling flow mass to fuel a powerful AGN, it is easy to see how radio galaxy activity can result from a cooling flow.  However, the impact of a radio galaxy on its environment is much less well understood.  In this paper, we use hydrodynamic simulations to study the evolution of a radio galaxy situated at the center of a galaxy/cluster.  In particular, we follow the evolution of the source at times after the jets have terminated.  We examine the interaction of the radio galaxy with the ISM/ICM and assess the environmental impact of the AGN activity.  Section~2 presents our set of hydrodynamic simulations and focuses on some numerical issues.  In Section~3, we briefly discuss the `active' phase of the source.  Section~4 presents our results for dead sources (i.e., the passive phase after the jets have turned off), and Section~5 discusses some astrophysical implications as well as the limitations of our calculations.  Section~6 draws our conclusions together. ",
        "conclusions": "Radio galaxies are known to be dynamical systems that evolve on timescales of $10^7-10^8\\yr$.  Hence, we expect many galaxies and clusters of galaxies to be host to dead, or relic, radio sources.  It is possible, and maybe even probable, that this past radio activity has influenced the energetics and thermodynamics of the hot, space-filling medium in these hosts.  Thus, a study of the environmental impact of dead radio galaxies is important and, given the high quality X-ray data now coming from the {\\it Chandra X-ray Observatory} and {\\it XMM--Newton}, very timely.  We have used high resolution hydrodynamic simulations to investigate one particular scenario relevant to dying radio galaxies --- i.e., a rather powerful radio-loud AGN situated at the center of a galaxy or cluster of galaxies whose activity ceases abruptly.  In particular, we have simulated back-to-back jets propagating in a $\\beta$-model galaxy/cluster atmosphere. We then shut down the jet activity and let the resulting structure evolve. To make the problem more tractable, we assume axisymmetry about the jet axis (thereby reducing the problem to two spatial dimensions) and neglect the action of magnetic fields.  This is a significant extension of the work of \\C00 and \\K00 since we simulate the active phase of the source and then shut off the jets, rather than modelling the late stages of evolution by simply letting a static bubble evolve under the action of buoyancy.  In the early lifetime of the simulated source (while the radio jets are still active), the shocked jet material forms a cocoon which is bounded by a shell of shocked ICM.  Even during the active phase, hydrodynamic instabilities start to shred this cocoon.  However, only after the jet activity ceases do KH and RT instabilities destroy the integrity of the cocoon.  Thereafter, the old cocoon material forms two plumes which rise in the cluster potential through the action of buoyancy forces.  These plumes entrain a significant amount of cooler material from the ICM core and lift this material high up into the cluster atmosphere.  At very late times, the galaxy/cluster atmosphere settles back into hydrostatic equilibrium but with a core specific entropy that has been enhanced by $\\sim 20\\%$ over its initial value. This entropy enhancement is due to the shocking of the lowest entropy material by the strong shock which bounds the radio cocoon during the early active phase.  We find that a large fraction ($\\sim 0.5$) of the injected energy is thermalized in the ISM/ICM.  Comparing late times to the initial state, we find that the ISM/ICM atmosphere has become inflated in order to maintain hydrostatic equilibrium after the thermalization of the jets energy.  Such a large thermalization efficiency raises the possibility that radio galaxies are important in the overall energy budget of the ISM/ICM."
    },
    "0201/astro-ph0201053_arXiv.txt": {
        "abstract": "{An automated search for star clusters close to the Galactic plane ($|b| < 5^o$) was carried out on the Point Source Catalogue of the DENIS survey. 44\\% of the Galactic plane have been observed and calibrated. The method allowed to retrieve 22 known star clusters and to identify two new ones, not published yet although previously presented in the 2MASS web site as embedded clusters in HII regions. Extinction in the field and in front of the clusters are estimated using a model of population synthesis. We present the method and give the properties of these clusters. ",
        "introduction": "The highly concentrated dust in the plane of the Milky Way makes difficult the observation in the visible range of embedded objects and most distant objects hidden behind the Galactic plane. However, the extinction decreases at longer wavelengths: in the K-band, it is about ten times smaller than in the visible. Thus the recent Near-Infrared surveys provide suitable data for systematic search for new objects projected on the Galactic plane. Using the \\textit{Two Micron All Sky Survey} \\citep[hereafter 2MASS,][]{Skrutskie1997} in the J (1.25 $\\mu m$), H (1.65 $\\mu m$) and K$_s$ (2.17 $\\mu m$) bands, \\cite{Dutra2000} carried out a systematic search for new star clusters in a 5$^o \\times 5^o$ field centered close to the Galactic plane and listed 58 IR star clusters or candidates. They also investigated selected regions with evidence of star formation and found 42 star clusters or candidates \\citep{Dutra2001}. \\cite{Hurt2000} serendipitously discovered two globular clusters near the Galactic plane. \\cite{Vauglin2001} undertook a search for galaxies by eye on J and K$_s$-band images from the \\textit{Deep-Near Infrared Survey of the Southern Sky} \\citep[hereafter DENIS,][]{Epchtein1997,Epchtein1999} and found 13 star clusters (Rousseau, private communication).  Using the data in I (0.82 $\\mu m$), J, and mostly K$_s$ bands from DENIS, we performed an automated search for star clusters in a $\\pm$5$^o$ band around the Galactic plane. In Sect.~\\ref{method} we describe our method and detection criteria. The basic properties of the new star clusters are given in Sect.~\\ref{results}. ",
        "conclusions": "We elaborated a program to perform an automated search of star clusters from the catalogues of extracted sources from the DENIS survey. We concentrated on the band around the Galactic plane ($|b| < 5^o$), where 44\\% of the data have already been calibrated. Most of the already known star clusters visible on the DENIS images in the probed region have been detected. Two star clusters, not published yet, have been identified. Both clusters are embedded in their associated HII regions. Extinction in front of the clusters is estimated. Uncertainties remain large on this determination mainly due to questionnable membership and to the photometric errors at faint magnitudes. Deeper and more accurate photometry and proper motions would allow to assert the cluster characteristics.  Although we have missed known star clusters and probably not yet detected clusters, our method applied to Near-Infrared data allows to search for old star clusters and embedded star clusters in a systematic way, with a lower efficiency than when looking at the image, but much more rapidly. We plan to run again this program once the complete region around the Galactic plane will be calibrated."
    },
    "0201/astro-ph0201265_arXiv.txt": {
        "abstract": "s{ Gravitational waves transport very detailed information on the structure and evolution of astrophysical sources. For instance a binary system in the early stages of its evolution emits a  wavetrain at specific frequencies that depend on the characteristics of the obital motion; as the orbit shrinks and circularize, due to radiation reaction effects, the orbiting bodies get closer and tidally interact. This interaction may result in the excitation of the proper modes of oscillation of the stars, and in the emission of gravitational signals that carry information on the mode frequencies, and consequently   on the equation of state in the stellar interior. These phenomena may occur either in solar type stars with orbiting planets and in compact binaries, and in this lecture we will discuss different approaches that can be used to study these processes in the framework of General Relativity. }  In recent years,  major progresses have been done in the construction of large ground-based interferometers, which will make accessible to observations a large frequency region, ranging from a few Hz to a few kHZ: TAMA has already performed some  observational runs, reaching a sensitivity that would enable to detect the coalescence of a neutron star-neutron (NS-NS) binary system occuring in our Galaxy with a signal to noise ratio higher than 30, and VIRGO, LIGO  and GEO600, are almost completely assembled and within the end of the year 2002 are expected to start scientific runs  \\cite{web}. Resonant bar detectors, ALLEGRO, AURIGA, EXPLORER, NAUTILUS, NIOBE, have been operational since many years, and are sensitive to a small frequency region of about 10-20 Hz, centered at $\\sim$1 kHz. As an example, Explorer and Nautilus would be able to detect, with signal to noise ratio equal to 1,  a burst of gravitational waves  of amplitude $h\\sim 4\\cdot 10^{-19}$, which would correspond to a mass of a few units in $10^{-3}~M_\\odot$ trasforming into gravitational waves at the center of our Galaxy \\cite{astone}. And finally, the space-based intereferometer LISA, which is expected to fly in about a decade, will enlarge the observational window to the low frequency region $10^{-4}~Hz\\lappreq \\nu \\lappreq 10^{-1}~Hz$.  Many are the astrophysical sources that, according to the theory of General Relativity, emit gravitational waves in the frequency region spanned by these detectors, and in this review I will discuss the characteristic properties of the signals emitted by some of the more interesting sources, and  show how to compute them. The detectability of a signal depends also on the {\\it event rate}, namely on how many events  per year occur in the volume accessible to observations, and  on the {\\it detection rate}, which depends on the responce of each detector to a specific source; how these rates are determined is,  of course, a very important issue, but a discussion of these problems is beyond the scope of this review, which will be focussed on the information that the gravitational signals carry on the structure (either geometrical or internal) of the source, on its motion, and on physical processes that may occur in some particular situations, like the resonant excitation of the stellar modes.  In Sec. I, I will shortly describe the quadrupole formalism, which is the easiest method to  estimate  gravitational fluxes and waveforms, and I will discuss its domain of applicability. In Sec. II and III, the quadrupole formalism will be applied to rotating neutron stars and to binary systems far from coalescence, respectively. In Sec. IV the resonant excitation of the g-modes in extrasolar planetary systems will be considered in the framework of a perturbative approach, and in Sec. V the gravitational emission of coalescing binary systems  will be studied. ",
        "introduction": " ",
        "conclusions": "\\label{sec6} In this review we have shown how to estimate the characteristic  features of the gravitational signals emitted by some  interesting astrophysical sources.  Chances to detect these waves crucially  depend  on the detailed knowledge of the waveform; indeed the standard technique used to extract  a non-continuous  signal from the noisy data of a detector is  the matched filtering technique, whose performances are very sensitive to a mismatching of the parameters: for this reason the templates must be as accurate as  possible. Since the most interesting signals are emitted in spacetime regions where the non linearity of the gravitational interaction plays an important role, the ultimate template would be that obtained by integrating the fully non linear equations of General Relativity. However,  fully nonlinear simulations of phenomena like the gravitational collapse or the coalescence of binary systems are, at present, at a preliminary stage, although major progresses in the field are underway. Thus, for the time being we need to rely on the results obtained  by approximate methods, like the quadrupole formalisms and the perturbative approach, to mention  those we have described in this paper. There are many things that we learn from these approaches. Consider for instance the evolution of a binary system: the quadrupole formalism, applied to point masses, allows to determine with high accuracy  the gravitational signal emitted when the two bodies are far apart, the frequencies and the expected amplitudes. Including radiation reaction effects and using the adiabatic approximation, we can predict how the orbit evolves, and how much time it takes to become circular.  However, if the system is composed by a solar type star and a close planet (or a brown dwarf companion), new phenomena may occur, like the resonant excitation of the g-modes of the star, which can be accounted for  only if  the  internal structure of the star is included in the picture; in this case, the perturbative approach provides the tools needed to evaluate how much energy is  radiated in gravitational waves because of the stellar oscillations,  at which frequencies, and for howlong can a system be in a resonant condition.  If the binary system is formed by very compact objects, like neutron stars or  black holes, the quadrupole approach works up to much higher frequencies, since in this case the masses  can really be treated, to a large extent, as point masses. During the latest phases of the inspiralling,  let say during the 10-20 cicles that preceed the merging when the signal frequency is of the order of 900-1000 Hz, the non linearity of the interaction produces deviations from the quadrupole signal which can be estimated by introducing post-newtonian corrections at  higher order of $(v/c)$ (see e.g. \\cite{damouriyersathia} and references quoted there). In the standard post-newtonian (PN) approach, the coalescing bodies are still considered as point masses, and this is good enough for black holes. However, in the case of NS-binaries the fundamental mode of the stars may be excited in the same region, and this excitation may produce corrections to the quadrupole signal that are comparable or higher than the PN corrections, and depend on the internal structure of the stars;  as we have shown in Sec. \\ref{sec5}, these phenomena can be studied again by a perturbative approach, which has been shown to work remarkably well, even when the two stars are at a distance as short as  three stellar radii \\cite{tutti2}.  An accurate detection of the signal emitted during the last few  cicles before coalescence may enlighten unresolved issues related to the internal composition of neutron stars and to the state of matter at supernuclear density. These effects, however will not be seen by the first generation of interferometers, for which the most likely source to be detected is the coalescence of BH-BH binaries with masses of the order of $20-30 ~M_\\odot$ \\cite{grishuck}. But there is hope for the future: advanced detectors are beeing proposed like EURO, a european project now under a feasibility study, that should be extremely sensitive at high frequency,  and would allow to study the processes we have described."
    },
    "0201/astro-ph0201115_arXiv.txt": {
        "abstract": "High-resolution $N$-body simulations of hierarchical cosmologies have shown that the density and velocity dispersion profiles of dark-matter haloes display well-definite universal forms whose origin remains unknown. In the present paper, we calculate the internal structure of haloes expected to arise in any such cosmologies by simply taking into account that halo growth proceeds through an alternate sequence of discrete major mergers and long periods of gentle accretion. Major mergers cause the violent relaxation of the system subject to the boundary conditions imposed by accreting layers beginning to fall in at that moment. Accretion makes the system develop inside-out from the previous seed according to the spherical infall model. The predicted structure is in very good agreement with the results of numerical simulations, particularly for moderate and low mass haloes. We find strong indications that the slight departure observed in more massive systems is not due to a poorer theoretical prediction, but to the more marked effects of the limited resolution used in the simulations on the empirical profiles. This may have important consequences on the reported universality of halo structure. ",
        "introduction": "The dominant dark component of matter in the universe appears to be clustered in bound haloes which form the skeleton of all astronomical objects of cosmological interest, from dwarf galaxies to rich galaxy clusters. All we know about these systems comes from their gravitational effects on the luminous matter they trap or on the light traveling across them. This is so little information that the following fundamental questions have prevailed for a long time. Does the internal structure of haloes depend on their mass? Does it depend on their past history? And on cosmology? How is this structure set?  Gunn \\& Gott (1972; \\citealt{gott75}; \\citealt{gunn77}) were the first to address this problem by considering the simplified case of haloes forming through spherical infall, that is, the collapse of a density fluctuation of smooth, isotropically distributed, dissipationless matter in an otherwise homogeneous expanding universe. Under the adiabatic invariant collapse approximation, they derived the density profile arising from some specific initial conditions. The effects of refining the derivation used and of adopting more and more realistic initial conditions have been subsequently addressed in a series of works (\\citealt{fg84}; \\citealt{berts85}; \\citealt{hs85}; \\citealt{rg87}; \\citealt{ryd88}; \\citealt{zh93}; \\citealt{lok00}; \\citealt{LH00}; \\citealt{pgrs}; \\citealt{ekp}; \\citealt{nuss}).  An important result of this research line is that haloes grow inside-out as new material is being incorporated in secondary infall. That is, despite the continuous shell-crossing (and radial relaxation, see below) of the infalling layers with the previously relaxed body, the structure remains at any moment essentially unaltered within the instantaneous radius. As shown below, this sole consideration should permit us to determine the halo density profile provided the infall rate of matter is known. Much progress has been achieved in the last twenty years in the modelling of halo growth. The extended Press-Schechter (PS) model (\\citealt{PS}; \\citealt{b}; \\citealt{BCEK}; \\citealt{LC93}) makes quite accurate predictions, indeed, on the rate at which haloes increase their mass \\citep{LC94} in hierarchical cosmologies like the one describing the real universe. Unfortunately, haloes develop in such cosmologies through continuous mergers rather than through smooth spherical infall.  The effects of mergers depend on the relative mass of the progenitors. Major mergers bring the whole system out of equilibrium making it move in the phase space around some attractor. During this process, particles experience random accelerations owing to the rapidly varying collective potential well which causes the relaxation of the system \\citep{LB67} in a similar although more dramatic way than in the case of spherically infalling layers. This violent relaxation erases any imprint of the previous history of the system, the new equilibrium state reached being characterised by a normal distribution of particle velocities similar to that yielded by the common two-body relaxation, although independent of particle mass. Were haloes isolated and, hence, unperturbed after any such dramatic event they would end up as infinite, spherical systems with a uniform, monomass, isotropic velocity dispersion. This is the reason why haloes are often modelled as isotropic, isothermal spheres (e.g., \\citealt{King72}; \\citealt{SIR99}). But haloes are not isolated systems. During the violent relaxation process, they keep on collecting matter through minor mergers making the boundary conditions to vary in some unknown way. This severely limits the predictive power of the violent relaxation theory (e.g., \\citealt{s87}).  To gain a deeper insight on the internal structure of haloes many authors have turned to numerical experiments. Using high-resolution cosmological $N$-body simulations Navarro, Frenk \\& White (1996, 1997) have found that the spherically averaged density profile of haloes of all masses is well fitted by the simple analytical expression \\begin{equation} \\rho(r)={\\roc \\rs^3 \\over r(\\rs + r)^2} \\label{nfw} \\end{equation} in any hierarchical cosmology analysed. In equation (\\ref{nfw}), $r$ is the radial distance to the halo centre, and $\\rs$ and $\\roc$ are the halo scale radius and characteristic density, respectively. The latter two parameters are related to each other and to the mass $M$ of the halo through the condition that the virial radius $R$ of the system encloses, by definition, an average density equal to a fixed factor $f$ times some density $\\rou$ characterising the universe at that moment. The values of $f$ found in the literature are in the range [178, 500], while $\\rou$ is taken equal to $\\rocr$, the critical density for closure, or $\\bar\\rho$, the mean cosmic density. Navarro and collaborators adopted $f=200$ and $\\rou=\\rocr$.  On the other hand, the spherically averaged, locally isotropised, velocity dispersion profile $\\Sigma(r)$ is found to be well fitted by the solution of the Jeans equation for hydrostatic equilibrium and negligible rotation, \\begin{equation} \\Sigma^2(r)\\left( {\\der \\ln \\Sigma^2 \\over \\der \\ln r} + {\\der \\ln \\rho \\over \\der \\ln r} \\right) = - {3\\,G M(r) \\over r} \\label{heq} \\end{equation} with $\\rho(r)$ the function given in equation (\\ref{nfw}) and $M(r)$ the corresponding mass within $r$, by adopting the boundary condition of null pressure at infinity. Note that, although the hydrostatic condition is natural to hold, the boundary condition at infinity is not obvious owing to the limited extent of haloes.  These empirical results have been confirmed by other authors (\\citealt{CL97}; \\citealt{HJS99a}; \\citealt{Bull01}), there being some controversy only at very small radii (\\citealt{Moore98}; \\citealt{JS00}; \\citealt{FM01}) where they are the most affected by the spatial resolution of the simulations. In any event, the fundamental question about the origin of that empirical halo structure remains. It is not even clear whether this is a general result of gravitational collapse, including smooth spherical infall, or the specific consequence of repeated mergers (\\citealt{HJS99b}; \\citealt{SW98}; \\citealt{Moore99}; \\citealt{LH00}).  In the present paper, we use a variant of the extended PS model distinguishing between minor and major mergers (\\S\\ \\ref{merg}) to derive the structure of haloes predicted in hierarchical cosmologies and compare it with that found in numerical simulations (\\S\\ \\ref{theo}). The result of this comparison supports the idea that the internal structure of haloes is the natural imprint of their hierarchical growth (\\S\\ \\ref{diss}). ",
        "conclusions": "\\label{diss}  It is generally believed that the PS formalism describes the mass growth of dark-matter haloes in hierarchical cosmologies but does not tell us anything about their internal structure. This is however in contradistinction with the idea that such a structure arises just from gravitation (a scale-free force) and {\\it the way that dark-matter haloes grow}. In the present paper, we have shown that, when the distinct dynamical effects of minor and major mergers are taken into account, the PS model also makes definite predictions on the internal structure of haloes, the resulting average theoretical profile being in very good agreement with that drawn from high-resolution numerical simulations. Our results therefore prove that the structure of haloes, including their scaling with mass, is the natural consequence of the combined action of minor and major mergers: outside the total radius of the system at formation the structure is fully determined by the rate at which mass is accreted, while, inside that radius, it results from the initial violent relaxation with boundary conditions set by accretion at that moment. The behaviour of the theoretical profiles derived here strongly suggests that the larger the halo mass, the more apparent are the effects on the respective density profiles of the limited resolution of the simulations. The confirmation of this suspicion would imply a different density profile for very massive haloes from that reported by Navarro, Frenk \\& White (1996, 1997) and, what is more important, the non-universality of halo structure.  The aim of the present paper was not to derive the accurate distribution of matter in real haloes, but to understand the origin of their empirical apparently universal structure. For this reason we have considered the simplified case of pure dark-matter haloes as those dealt with in \\nbody experiments. In the real universe, about one tenth of the halo mass is in the dissipative baryonic component, which might have appreciable effects. Likewise, we have assumed spherical symmetry and neglected any rotation of haloes as well as any anisotropy of the velocity tensor while, in hierarchical cosmologies, haloes have a slight angular momentum and, what is more important, they are immersed in large filamentary structures making them accrete matter preferentially along one privileged direction \\citep{West94} and feel, in their final steady configuration, the tidal field of such anisotropic structures \\citep{ss93}. On the other hand, even in the case of exact spherical symmetry, some velocity anisotropy would emerge in the outer part of haloes owing to the distinct evolution of the radial and tangential velocity dispersions in accreted layers during infall. All these secondary effects explain the elongation and slight angular momentum and velocity anisotropy observed in real as well as simulated haloes. (Note however that the density profile derived here is independent of the actual degree of anisotropy of the velocity tensor.) Finally, in the present paper, we have focused on the structure of haloes at $z=0$, the only redshift for which accurate empirical data for different cosmologies are available (\\citealt{NFW97}). In a forthcoming paper, we will study the predicted dependence on redshift of this structure and compare it with the results of numerical simulations carried out by \\citet{Bull01} for a $\\Lambda$CDM cosmology.  \\vspace{0.75cm} \\par\\noindent {\\bf ACKNOWLEDGMENTS} \\par  \\noindent This work was supported by the Spanish DGES grant AYA2000-0951. AM acknowledges the hospitality of the CIDA staff in M\\'erida (Venezuela) where part of this work was carried out."
    },
    "0201/astro-ph0201323_arXiv.txt": {
        "abstract": "{New colour distributions have been derived from wide field UBVRI frames for 36 northern bright elliptical galaxies and a few lenticulars. The classical linear representations of colours against $\\log r$ were derived, with some improvements in the accuracy of the zero point colours and of the gradients. The radial range of significant measurements was enlarged both towards the galaxy center and towards the outskirts of each object. Thus, the ''central colours\", integrated within a radius of 3 \\arcsec, and the ''outermost colours\" averaged near the $\\mu_V=24$ surface brightness could also be obtained. Some typical deviations of colour profiles from linearity are described. Colour-colour relations of interest are presented. Very tight correlations are found between the U$-$V colour and the $Mg_2$ line-index, measured either at the galaxian center or at the effective radius.} ",
        "introduction": "The ''classical\" data on the large scale colour distributions of E-type galaxies relies on observations by Bender and M\\\"ollenhof (1987), Vigroux et al. (1988), Franx et al. (1989), Peletier et al. (1990), Goudfrooij et al. (1994), to quote only the papers discussing the 1-D profiles of colour against radius, as distinguished from studies of dust patterns. Most of these data were reconsidered by Michard (2000) (RM00), in an attempt to collect a significant sample of objects with a complete optical colour set, i.e. U$-$B, B$-$V, B$-$R and V$-$I in a coherent photometric system. This was adequate to confirm previous indications about the cause of colour gradients: these appear to be due essentially to population gradients within galaxies, with the dust playing no important role, except in galaxies with central intense dust patterns. Such objects are rather rare among the Es.  Similar to most spectral indices of stellar populations, the colours suffer from the well known age-metallicity degeneracy, and, except U$-$B or U$-$V, are not very sensitive to the two parameters. They are affected by dust, at least locally, or perhaps systematically in the central regions according to inferences based on a survey by Michard (1999) (RM99). On the other hand, they may be measured at lower surface brightnesses or larger radii than the line indices. They could therefore bring useful information to the study of fossil stellar populations, and further constraints upon models of the evolution of E galaxies. The present work aims to provide an enlarged sample of objects with complete colour data, extending farther in radius than in previous studies, and hopefully of improved accuracy.  In Paper I, we present the usual information about the observations and data reduction, and part of the results in tabular form. A larger set of results will be made available in electronic form. The frames, partly reduced, will be made available from the HYPERCAT database, Observatoire de Lyon.  In Paper II, under the assumption that the observed colour gradients reflect abundance variations along the radius, metallicity gradients will be computed from the present data, using new colour-metallicity calibrations derived from multi-population models for E-galaxies. These metallicity gradients allow an estimation of central and mean metallicities. Statistics of galaxies included in our sample indicate that mean metallicities are about solar, in agreement with the study by Trager et al. (2000) based on spectral indices.  \\vspace{0.5 cm}  {\\em Often used notations} \\begin{itemize} \\item $r$ isophotal radius; $r=(ab)^{1/2}$ for an ellipse of semi-axis $a$ and $b$. \\item $\\Delta_{UB}$ colour gradient in U$-$B; $\\Delta_{UB}=\\mathrm{d(U-B)/d}(\\log r)$ and similar for other colours. \\item diE, boE, unE: subclassification of E galaxies as disky, boxy and undetermined. \\end{itemize} ",
        "conclusions": "New colour-radius relations have been derived for 36 E-type galaxies of the northern Local Supercluster, using UBVRI frames obtained with the 120cm telescope of the Observatoire de Haute-Provence. Four SA0, i.e. NGC3115, 3607, 4551 and 5866 were also observed.  We aimed to take advantage of the large field of the camera to observe the galaxies at larger radii than hitherto feasible, and thus improve the accuracy of colour gradients. The availability of the series of aperture photometry in PP88 and PN94 for most of the sample was also considered an asset towards a more coherent system of colours. It appears indeed that {\\em the colour calibrations are improved here compared to previous work}, if this can be judged from the quality of correlations between zero point colours in various surveys (see Sect. 4.1)  Two steps in the reduction procedure were thought significant in improving the quality of colour profiles: the first was the adjustment of the FWHM of the PSFs in a given colour set of 5 frames to the best of the five. This allowed us to get significant colours much closer to the galaxy center than otherwise feasible. The second was a careful ''mapping\" of the background of each frame, in order to lessen the background fluctuations remaining after the usual flat-fielding procedures. Both these precautions proved successful, and, as a result, the radial range of satisfactory colour measurements was greatly enlarged. Near the galaxy center, it proved feasible to obtain ''central colours'', i.e. colours integrated in the circle $r=3\\arcsec$, in fair agreement with high resolution data (see Sect. 3.2.7.2 and Table 10).  On the other hand, colours could be obtained at much lower surface brightness (or larger radii) than in previous work. Our colour data extend to $23.2 <\\mu_V <25.2$, with a median value near $\\mu_V=24.5$ in all colours. According to the comparisons in Sect. 4.1, this is {\\em 1.5 to 2 magnitudes deeper than in previous work}. ''External colours\", refering to the level  $\\mu_V=24.5$ whenever possible, are published for the first time (see Table 11), and may be useful to give some indications about stellar populations at the outskirts of E-galaxies.  On the other hand, the ''red halo\" effect of the camera was found to give enormous errors in V$-$I colours and gradients. These were corrected by a rigorous technique, and results in agreement with ''classical\" data were obtained. Considering the V$-$I gradients, one is not happy however to introduce in their evaluation, corrections larger than the quantity to be measured! Besides this specific problem of the red halo of thinned CCD, the far wings of the PSFs have been proven in a recent paper (RM01) to have non negligible effects in the gradients of other colours, and also to vary with the age of mirror coatings. It is not impossible that the U$-$B or U$-$V gradients given here are overestimated by 15-20\\%, although they are in excellent statistical agreement with the well-known work of Peletier et al. (1990).  Various colour gradients against $\\log r$ for a given object are well correlated, generally better than in previous work (see statistics in Sect. 4.1), which is interpreted as due to smaller measuring errors, notably in U$-$B. These improvements in accuracy did not bring out any obvious correlation between gradients and other galaxy properties. A few galaxies have exceptionally steep colour gradients (nearly at $2\\sigma$) without sharing other properties.  Colour-colour relations can be built from the present data for several locations in galaxies, such as near center, various fractions of the effective radius $r_e$, or the ''outermost\" measured range around $\\mu_V=24$. All these diagrams overlap to form a single stripe with moderate scatter (except for one rather obvious calibration error?). These might prove useful to test theories of old stellar populations and of their host galaxies. Colour-colour diagrams based upon integrated colours have already been used for this purpose (Worthey, 1994).  The U$-$B or U$-$V colours correlate very well with the $Mg_2$ index, both near the galaxy center and at the effective radius $r_e$. This seems to rule out any large influence of diffuse dust in the colours and colour gradients in E-galaxies. This was considered likely by Witt et al. (1992) and discussed by Wise and Silva (1996) with inconclusive results. Previous arguments against such an influence were presented in RM00: they were based upon the relative average values of the gradients in various colours, and are reinforced in the present work, since the mean gradients are nearly unchanged, and their errors lessened.  It is well known that, for single-burst stellar populations, colours and line indices depend both on the metallicity and on the age of the system (Worthey 1994; Borges et al. 1995). However, E-galaxies are constituted by a population mix, having age and metallicity distributions which reflect their star formation histories. Therefore a colour-metallicity calibration requires the use of models able to provide those distributions and, consequently, the integrated colours along the galaxy lifetime. Such a calibration will be presented in Paper II."
    },
    "0201/astro-ph0201445_arXiv.txt": {
        "abstract": "Deep images in the 10 $\\mu$m spectral region have been obtained for five massive Galactic globular clusters, NGC~104 (=47~Tuc), NGC~362, NGC~5139 (=$\\omega$~Cen), NGC~6388, NGC~7078 (=M15) and NGC~6715 (=M54) in the Sagittarius Dwarf Spheroidal using ISOCAM in 1997.  A significant sample of bright giants have an ISOCAM counterpart but only $<$ 20\\% of these have a strong mid-IR excess indicative of dusty circumstellar envelopes.  From a combined physical and statistical analysis we derive mass loss rates and frequency. We find that {\\it i)}~significant mass loss occurs only at the very end of the Red Giant Branch evolutionary stage and is episodic, {\\it ii)} ~the modulation timescales must be greater than a few decades and less than a million years, and {\\it iii)}~mass loss occurrence does not show a crucial dependence on the cluster metallicity. ",
        "introduction": "A complete, quantitative understanding of the physics of mass loss processes and the precise knowledge of the gas and dust content in Galactic globular clusters is crucial in the study of Population II stellar systems and their impact on the Galaxy evolution. It also has major astrophysical implications on related problems such as the ultraviolet excess found in elliptical galaxies (Greggio \\& Renzini 1990; Dorman et al. 1995) and the interaction between the intracluster medium and the hot halo gas (cf. e.g. Faulkner \\& Smith, 1991).  Indirect evidence for mass loss in Population II stars includes the observed morphology of the horizontal branch (HB) in the cluster color--magnitude diagrams, the pulsational properties of the RR Lyrae stars, and the absence of asymptotic giant branch (AGB) stars brighter than the red giant branch (RGB) tip (Rood 1973, Fusi Pecci \\& Renzini 1975; 1976, Renzini 1977, Fusi Pecci et al. 1993, D'Cruz et al. 1996). The expected mass loss is about $0.2\\,M_{\\odot}$ along the RGB and about $0.1\\,M_{\\odot}$ along the AGB (e.g. Fusi Pecci \\& Renzini 1976). As a consequence of such mass loss processes, dust and gas should be present in the intracluster medium, as diffuse clouds (cf. e.g. Angeletti et al. 1982) or concentrated in circumstellar envelopes. If no cleaning mechanism is at work between two Galactic plane crossings, a few tens solar masses of intracluster matter should be accumulated in the central regions of the most massive clusters (i.e. those with large central escape velocity).  Early searches (see Roberts 1988, Smith et al. 1990 and references therein) for this intracluster gas (HI, HII, CO) and more recently detection of ionized gas in 47 Tuc (Freire et al. 2001), yielded upper limits or marginal detections well below 1~$M_{\\odot}$.  Mid--IR excesses and scattered polarized light have also been observed in the central region of massive globular clusters (Frogel \\& Elias, 1988; Gillet et al. 1988; Forte \\& Mendez 1989; Minniti et al., 1992; Origlia, Ferraro \\& Fusi Pecci 1995, Origlia et al. 1997b).  They are mainly associated with long period variables evolving along the AGB.  These preliminary results seem to indicate that in globular clusters even the brightest IR point sources are fainter than 1 Jy in the 10 $\\mu$m spectral region and most of the emission comes from warm circumstellar dust with a minor photospheric contribution. Any diffuse component is much fainter than 10 mJy$\\,{\\rm arcsec^{-2}}$.  Other evidence of dust in the intracluster medium of globular clusters comes from the presence of IRAS sources (Lynch \\& Rossano 1990, Knapp \\& Gunn 1995, Origlia et al. 1996) in their central regions and from more recent surveys using ISOPHOT (Hopwood et al. 1999) and ground based sub-millimeter antennas (cf. e.g. Origlia et al. 1997a, Penny, Evans \\& Odenkirchen 1997, Hopwood et al. 1998).  However, all these studies show that globular clusters are deficient in diffuse intracluster matter and that some mechanism(s) must be at work to remove the gas and dust which should have accumulated between each passage through the Galactic plane (cf. e.g. Faulkner \\& Smith, 1991).  Given this observational scenario the desirability of a deep mid-IR survey of the central regions of globular clusters in the continuum and in selected dusty features is obvious.  Such a deep, spatially resolved photometric survey became possible with the spectrophotometric capabilities of ISOCAM on board of the Infrared Space Observatory (ISO, Kessler et al. 1996). ISOCAM (Cesarsky et al. 1996) provided relatively fine spatial resolution, large field coverage, and high sensitivity in the 10 and 20 $\\mu$m spectral regions.  Ramdani \\& Jorissen (2001) performed a deep survey at 12 $\\mu$m in the external regions of 47~Tuc at different distances from the cluster center, with the specific goal of studying the mass loss during the AGB evolutionary stage.  They cross-correlated their ISOCAM survey with the DENIS survey and found dust excess only in two well known bright variables, confirming the strong link between stellar pulsation activity and mass loss modulation.  With the goal of studying the mass loss during the RGB evolutionary stage, a deep survey of the very central regions of six, massive globular clusters has been made using ISOCAM with two different filters in the 10 $\\mu $m spectral region.  Mid-IR observations have the advantage of sampling an outflowing gas fairly far from the star (typically, tens/few hundreds stellar radii).  At 10 $\\mu$m the sampled material typically left the star a few decades previously. Thus mass loss rates inferred from the mid-IR will smooth the mass loss rate over a few decades. For astrophysical purposes it is the longterm average mass loss which is important. Conceptually one could sample different distances and smoothing times by observing at different wavelengths, however at this time that is not practical. At longer wavelengths detectors lack the requisite sensitivity; at shorter wavelengths the photosphere dominates. ",
        "conclusions": "Our ISOCAM survey of the central region of six massive globular clusters provided mid-IR photometry for 78 sources with identified stellar counterparts.  52 sources are associated with bright giants close to the RGB-tip. Of these about 40\\% have strong mid-IR excess ascribed to the presence of dusty circumstellar envelopes. Correcting for stars not detected because of the low spatial resolution of ISOCAM, dusty envelopes are inferred around about 15\\% of the brightest giants ($M_{\\rm bol}\\le-2.5$).  The inferred mass loss rates are in the range $10^{-7} < dM/dt < 10^{-6} M_{\\odot}\\,{\\rm yr^{-1}}$, assuming as reference values an outflow velocity of 14~${\\rm km\\,s^{-1}}$ and a gas-to-dust mass ratio of 200 at the metallicity of 47~Tuc and for a stellar luminosity of 1000~$L_{\\odot}$.  The major astrophysical implications are: \\begin{itemize} \\item The mass loss occurs very near the RGB-tip and is episodic.  \\item The mass loss episodes must last longer than a few decades and less than a million years.  \\item There is no indication for a strict metallicity dependence of the frequency of mass loss occurrence. \\end{itemize}"
    },
    "0201/astro-ph0201390_arXiv.txt": {
        "abstract": "\\noindent  We present infrared photometry of the Galactic Bulge X-ray binary systems GX13+1 and GX17+2 obtained in 1997 July and August using OSIRIS on the 1.8m Perkins Telescope at Lowell Observatory.  GX13+1 clearly varies over $\\sim$0.6 magnitudes in the $K$-band.  Our light curve suggests a modulation on a timescale of $\\sim$20 days, which is in agreement with previously proposed orbital periods for the system.  The IR counterpart of GX17+2 is also variable in the $K$-band over $\\sim$0.8 magnitudes on a timescale of days to weeks, extending the variability first seen by Naylor, Charles, \\& Longmore (1991).  We discuss the implications our data have for Deutsch \\etal's (1999) identification of ``star A'' as the true IR counterpart of GX17+2.  The variability observed in our photometry of the blend of star A and the foreground star NP Ser implies a $\\sim$4 magnitude intrinsic variation in the $K$-band for GX17+2. ",
        "introduction": "In low-mass X-ray binaries (LMXBs), mass is transferred from a late-type star to its highly compact companion, either a neutron star or a black hole, via an accretion disc.  These systems can be classified according to location within the Galaxy, accretion characteristics, and luminosity (van Paradijs \\& McClintock 1995).  Mapping of the X-ray sky has revealed a dozen bright X-ray sources within $15^{\\circ}$ longitude and $2^{\\circ}$ latitude of the Galactic Centre (Warwick \\etal 1988).  Known as the ``galactic bulge'' or ``bright bulge'' sources (henceforth referred to as ``GBS''), these LMXBs are among the most luminous X-ray sources in the Galaxy (typical $L_{X} \\sim 10^{38}$ \\ergs).  However, the GBS remain the most poorly understood group of LMXBs, due to the heavy obscuration in the direction of the Galactic bulge which makes optical study nearly impossible.  The GBS are thought to be neutron star LMXBs, and quasi-periodic oscillations (QPOs) have been detected in several systems (\\eg van der Klis 1989).  Attempts to detect orbital variability in the GBS have generally been unsuccessful, suggesting that their periods may be longer than those of canonical LMXBs (Charles \\& Naylor 1992, hereafter CN92).  On the basis of their X-ray colour-colour diagrams, neutron star LMXBs have also been divided into two classes, known as $Z$ and atoll sources (see van der Klis 1995 for a review).  In this scheme, six of the brightest neutron star LMXBs (which includes several of the GBS) are classified as $Z$ sources, while the remainder fall into the atoll category.  However, recently this distinction has come into question, as several atoll sources show $Z$-type colour-colour diagrams when examined on a long timescale (Muno \\etal 2001, Gierlinski \\& Done 2001).  Nevertheless clear differences remain between the X-ray luminosity and spectral evolution of the canonical $Z$ sources and that of the atolls.   The infrared provides us with an ideal window for observing these highly obscured systems.  Observing in the IR has two primary advantages:  the late-type secondaries in LMXBs are brighter relative to the accretion discs, and, more importantly for the GBS, the ratio of $V$- to $K$-band extinction is nearly 10 (Naylor, Charles, \\& Longmore 1991, hereafter NCL91).  Over the past eight years, we have developed a program of IR observations of X-ray binaries (XRBs), beginning with the discovery via colours or variability of candidates for the IR counterparts to the X-ray sources using precise X-ray and radio locations (NCL91).  Following this photometric survey we performed a spectroscopic survey of LMXBs, obtaining IR spectra of a number of systems, including Sco X-1, GX13+1, and Sco X-2 (Bandyopadhyay \\etal 1997 and 1999, hereafter B97 and B99).  The spectra enabled us to place constraints on the spectral types of the mass donors in these systems, indicating that the secondaries in the GBS may be evolved rather than main-sequence stars.   If, as the evidence suggests, the companion stars in the GBS are indeed evolved stars, we would expect any variability in the IR light curve due to the orbital period to occur on a timescale of days or weeks, rather than hours as in canonical LMXBs.  To search for long orbital periods in the atoll LMXB GX13+1 and the $Z$ LMXB GX17+2 (as classified by Hasinger \\& van der Klis 1989), in 1997 we obtained IR photometry of both sources over a period of approximately 6 weeks.  An IR study of the GX13+1 field mapped by the {\\it Einstein HRI} in X-rays revealed a variable IR source at the Grindlay \\& Seaquist (1986) radio position of GX13+1 (CN92).  The observed variability made this IR source a very strong candidate for the counterpart to GX13+1.  The $K$-band spectrum we obtained in 1997 showed Brackett $\\gamma$ emission, confirming the identification of the IR counterpart (B97).  An additional spectrum obtained in 1999 with improved S/N exhibited both Brackett $\\gamma$ emission and CO absorption bands; using the technique of optimal subtraction, we determined the most likely spectral type of the secondary to be a K5 {\\sc iii} (B99).  In contrast to GX13+1, detection of the true IR counterpart to GX17+2 has proven more difficult.  GX17+2 was optically ``identified'' on positional grounds alone with a G star now known as NP Ser more than 25 years ago (Tarenghi \\& Reina 1972).  However, the absence of any optical variability or spectroscopic peculiarity in a star purported to be associated with one of the most luminous X-ray sources in the Galaxy has long been a major puzzle, with most explanations requiring that the G star be simply a line-of-sight object.  Furthermore, the IR counterpart to NP Ser has shown variability, and a consistent fit for the optical and IR colours, extinction, distance, and spectral type could not be found (NCL91).  The situation has recently been clarified by HST observations of GX17+2 which resolved NP Ser into two components, the fainter of which, labelled ``star A'' by Deutsch \\etal (1999), has been suggested as the true IR counterpart to GX17+2.  The unusual nature of star A has been demonstrated by Callanan \\etal's (1999) Keck photometry which suggests that star A exhibits extremely large ($\\sim$3.5-4 mag) IR variability.  In this paper we present long-term ($\\sim$5-week) light curves of GX13+1 and GX17+2 from July and August 1997; both sources show clear variability of greater than 0.5 magnitudes in the $K$-band. ",
        "conclusions": "Our $K$-band photometry of GX13+1 shows evidence for a $\\sim$20 d modulation.  We note that a 20 d period correlates extremely well with that expected for a Roche-lobe filling binary with a K5{\\sc iii} secondary (B99).  On the basis of the observed variability in the IR together with the identification of the secondary's spectral type, we suggest that the observed $\\sim$20 d modulation may be the orbital period of GX13+1.  We note that to conclusively identify any period of order $\\sim$20 d in GX13+1, it will be necessary to obtain a well-sampled IR light curve with a baseline of $\\gtsimeq$60 d.  We also find evidence in the GX13+1 ASM X-ray light curve for quasi-periodicity on a timescale of 20-30 days, but this modulation is not consistently present throughout the $\\sim$5-year dataset.  We also observe a $\\sim$10 d modulation in the IR light curve of GX17+2.  Our photometry was performed on the blend of the foreground star NP Ser and the recently identified counterpart candidate star A; the observed variability is consistent with that expected as a result of a large modulation of the fainter of the two stars.  The IR photometric modulation and revised position for GX17+2 strongly supports the identification of star A as the true counterpart to the X-ray source.  Previous attempts at obtaining an IR spectrum of GX17+2 have been unsuccessful, no doubt as a result of the small separation between NP Ser and star A as well as star A's large photometric variability which would make it nearly undetectable for a significant fraction of its IR period.  However, now that these difficulties are recognized, future IR spectroscopic observations of star A will hopefully prove more successful in investigating the physical nature of the extraordinary variability of the true IR counterpart."
    },
    "0201/astro-ph0201029_arXiv.txt": {
        "abstract": "{We explore the scale and angular dependence of the cosmic shear three-point correlation function. The shear field is found to have a much more complex three-point function than the convergence field, but it also exhibits specific motifs which draw signatures of gravitational lensing. Exact shear patterns are inferred analytically for some geometrical shear triplets configurations, when simple interpretations can be derived. {A more complete description of their geometrical properties has then been carried out} from ray-tracing numerical simulations and simple models. These patterns reveal  typical  features that can be utilized for non-gaussian signal detection in cosmic shear surveys. We test the robustness of these properties against complex noise statistics and non-trivial survey topologies. From these conclusive checks, we predict that the VIRMOS-DESCART survey should allow a detection of a non-gaussian signal with a comfortable significance for a low matter density Universe. ",
        "introduction": "The detections of cosmic shear signal (van Waerbeke et al. 2000, Bacon et al.  2000,  Wittman et al.   2000, Kaiser et  al. 2000, Maoli et al 2001) opened the analysis of  large-scale mass distribution in the Universe to weak lensing  surveys. As the surveys sizes progressively increase, the noise and systematics are reduced to  reasonable levels, which permits to constrain cosmological parameters from the amplitude and the shape of the shear two-point correlation function (Maoli et al. 2001, van Waerbeke et al. 2001b). However, although it is weakly sensitive to the cosmological constant $\\Omega_{\\Lambda}$ (Bernardeau, van Waerbeke \\& Mellier 1997),   the two-point function is a degenerate combination of the matter density $\\Omega_{{\\rm m}}$ of the Universe and the amplitude of the power spectrum $\\sigma_8$ (Villumsen 1996, Jain \\& Seljak 1997).  \\\\ This degeneracy can be broken from the angular dependence of the cosmic shear amplitude (Jain \\& Seljak 1997), but it relies on the prior knowledge of the shape of the mass power spectrum. An alternative is to directly probe weak shear maps,  which contain non-gaussian features that can be used to derive $\\Omega_{{\\rm m}}$ with the only assumption that initial conditions were gaussian (Bernardeau, van Waerbeke, Mellier 1997, van Waerbeke, Bernardeau, Mellier 1999) - this later assumption being eventually testable in the data set itself . \\\\ So far, all the theoretical predictions regarding non-gaussian features are based on a critical reconstruction process of either a convergence map (filtered for instance with a top-hat window function) or an aperture mass map (Kaiser et al. 1994, Schneider 1996, Schneider et al. 1998, Bernardeau \\& Valageas 2000).  Unfortunately, the panoramic reconstruction of mass maps from real data turned out to be considerably more difficult than expected.  The masking process discussed by van Waerbeke et al. 2000 produces patchy surveys with a non-trivial topology and inhomogeneous noise. The resulting mass maps have poorly understood  statistical properties which is practically difficult to handle with confidence. We therefore explore another option which uses direct signatures of gaussian effects in shear map patterns.  Shear pattern study is a priory more difficult because the third order moment of the local shear vanishes for obvious symmetry reasons. It is therefore necessary to seek for peculiar configurations (or geometries) of shear triplets for which a significant signal is expected. This is the aim of this paper. In the next section, we present the exact analytical results obtained for specific geometries, their interpretation and the results obtained from ray-tracing simulations which exhibit the complete three-point function patterns. The detection of these features in these simulations is presented in section 3. These results suggest a detection strategy in cosmic shear surveys which is put forward and tested with mock catalogs in section 4. We checked carefully that the masks do not compromise the possibility  of detecting the effect. Finally, we estimate the expected error level of such measurements in the available data sets. ",
        "conclusions": "We have presented various geometrical patterns that the shear three-point correlation function is expected to exhibit. Its dependence with the cosmological parameters is found to be similar to that of the skewness of the convergence field, opening an alternative way to break the degeneracy between the amplitude of the density fluctuations, $\\sigma_8$, and the density parameter of the Universe $\\Omega_{{\\rm m}}$. However, the reduced three-point function of the shear is more dependent on the power spectrum index than the skewness. In particular it is expected to vanish when the index gets close to $-2$. Numerical investigations have nonetheless proved the shear patterns to be robust enough to provide a solid ground for the detection of non-gaussian properties in cosmic shear fields.  We proposed a detection strategy that has been tested in mock catalogues that include realistic noise structures such as residual systematics and PSF anisotropy as seen in the real data. The quality of the PSF correction is always good enough to provide an accurate measurement of the shear three-point function. Since the mock catalogues were designed to reproduce the characteristics of the current VIRMOS-DESCART lensing survey, we conclude that non-gaussian signal should be detectable in this data set. The result of our investigations is presented in another paper (Bernardeau et al. 2002).  Beyond the detection, the scientific exploitation of the 3-point function for cosmology also depends on our ability to overcome other important difficulties. For instance, we found the cosmic variance amplitude to be of the order of 30\\% of the signal on the reduced three-point function, in agreement with previous studies made for the convergence. Other issues regarding source clustering, source redshift uncertainties, and intrinsic alignment of galaxies have not been considered here.  The measurement of non-gaussian signatures in lensing surveys is of course of great interest because it provides an independent measure of the mean mass density of the Universe in addition to test the gravitational instability paradigm which lead to large scale structures. It is likely that the analysis of cosmological non-gaussian signatures will be one of the major and most promising goals of emerging dedicated lensing surveys\\footnote{Such as the Canada-France-Hawaii Telescope Legacy Survey , http://www.cfht.hawaii.edu/Science/CFHLS/} that take advantage of panoramic CCD cameras of the MEGACAM generation (Boulade et al 2000).  {"
    },
    "0201/astro-ph0201244_arXiv.txt": {
        "abstract": "We have undertaken a long-term project, Planets in Stellar Clusters Extensive Search (PISCES), to search for transiting planets in open clusters. As our first target we have chosen NGC 6791 -- a very old, populous, metal rich cluster. In this paper we present the results of a test observing run at the FLWO 1.2 m telescope. Our primary goal is to demonstrate the feasibility of obtaining the accuracy required for planetary transit detection using image subtraction photometry on data collected with a 1 m class telescope. We present a catalog of 62 variable stars, 47 of them newly discovered, most with low amplitude variability. Among those there are several BY Dra type variables. We have also observed outbursts in the cataclysmic variables B7 and B8 (Kaluzny et al. 1997). ",
        "introduction": "} Since antiquity mankind has wondered whether planetary systems other than our own exist. The first documented effort aimed at extrasolar planet detection was undertaken by Huygens in the XVIIth century. Starting in the 1930s, subsequent searches have been attempted, but failed to produce positive results due to insufficient measurement precision.  Only recently it has become possible to obtain radial velocity measurements accurate enough to indirectly detect planets via Doppler shifts of stellar spectra of stars other than the Sun (Mayor \\& Queloz 1995).  To date, over 70 planets have been discovered, mainly around solar type stars. All of them were detected by radial velocity surveys (eg. Marcy et al.~2001, Noyes et al.~1997). For one of these systems, HD 209458, the transit of the planet across the host star's disk has been observed (Charbonneau et al.~2000, Henry et al. 2000), thus demonstrating the feasibility of detecting planets this way. Several groups are currently monitoring the brightness of thousands of stars to search for planets via transits (eg. Brown \\& Charbonneau 1999, Quirrenbach et al.~1998). Recently, Udalski et al.\\ (2002) discovered 46 stars with transiting low-luminosity companions. These objects may be planets, brown dwarfs or M dwarfs. A confirmation of their nature will be provided by mass determinations based on photometry combined with radial velocities from followup spectroscopic observations.  The analysis of the properties of stars with planets suggests that they are on the average significantly more metal rich than those without (Santos et al.~2001). Some studies indicate that the source of the metallicity is most likely ``primordial'' (Santos et al.~2001, Pinsonneault et al.~2001), while others suggest that the observed high metallicity is intrinsic only in some cases, with the more likely cause being the accretion of planetesimals onto the star (Murray \\& Chaboyer 2001) or the infall of giant gas planets (Lin 1997).  Although the observed lack of planets in the low metallicity ([Fe/H]$=-0.7$) globular cluster 47 Tuc (Gilliland et al.~2000) is compatible with the ``primordial'' metallicity scenario, the case is far from being resolved. In such dense environments as the cores of globular clusters encounters with other stars may lead to the breakup of planetary systems (Davies \\& Sigurdsson 2001). The study of open clusters offers the possibility of disentangling the effects of metallicity and crowding: their stellar densities are not high enough for crowding-induced disruption to be effective.  We have undertaken a long-term project, Planets in Stellar Clusters Extensive Search (PISCES), to search for transiting planets in open clusters. As our first target we have chosen NGC 6791 $[(\\alpha,\\delta)_{2000}=(19^h20.8^m, +37^{\\circ}51')]$, a very old, extremely metal rich cluster ($\\tau$=8 Gyr, [Fe/H]=+0.4; Chaboyer et al.~1999). At a distance modulus of (m-M)$_V$ = 13.42 (Chaboyer et al.~1999) it contains about 10000 stars (Kaluzny \\& Udalski 1992, hereafter KU92).  In this paper we present the results of a test observing run at the FLWO 1.2 m telescope. Our aim was to demonstrate the feasibility of obtaining the required accuracy using image subtraction photometry on data collected with a 1-m class telescope needed to reliably detect transits of inner-orbit gas-giant planets with an acceptably low false alarm rate.  On clear nights we have reached the desired level of photometric precision. Unfortunately, there were few such nights during our 26 night run, which is typical for July at FLWO. Even though this dataset is not optimal for planetary transit detection, it allowed us to discover 47 new low amplitude variables, compared to 22 previously known (Kaluzny \\& Rucinski 1993; Rucinski, Kaluzny \\& Hilditch 1996, hereafter: RKH).  The paper is arranged as follows: \\S 2 describes the observations, \\S 3 summarizes the reduction procedure, \\S 4 outlines the procedure of variable star selection, \\S 5 contains the variable star catalog. Concluding remarks are found in \\S 6. ",
        "conclusions": "} In this paper we have demonstrated the feasibility of obtaining photometry accurate enough to detect planets through transits in open clusters with 1 m class telescopes. The analysis of the data collected for this purpose resulted in the discovery of 47 new low amplitude variables, compared to 22 previously known (Kaluzny \\& Rucinski 1993, RKH).  The first stage of our project will be to obtain for this and several other clusters continuous observations with a 1-m class telescope for about $\\sim$30 nights under very good observing conditions. The photometry will be obtained in two bands, $R$ and $V$, to allow us to differentiate between planetary transits and blended eclipsing binaries. Planetary transits should be gray, as the planet does not contribute any measurable light to the system, while the superposition of an eclipsing binary with another star will show a change in color during the eclipse.  In the next stage better light curves will be obtained with a 2 m class telescope for selected planet transit candidates. Using radial velocity measurements derived from spectroscopic observations with a 6-10 m class telescope it will be possible to distinguish planetary and brown dwarf transits from grazing eclipses by main sequence companions. A precision of 1 km s$^{-1}$ will enable us to identify and reject stars with companions above 0.0075 M$_\\odot$ ($\\sim 8$ M$_{J}$)."
    },
    "0201/astro-ph0201072_arXiv.txt": {
        "abstract": "I review recent results from numerical simulations on the structure and dynamics of the ISM, and attempt to put together a coherent dynamical scenario. In particular, I discuss results on 1) the spatial distribution of the gas components, showing that reasonable agreement between simulations and observations exists, but noting that in most models the components are simply defined as temperature intervals, because distinct thermodynamic ``phases'' do not arise; 2) some statistical issues of the physical fields, like the dependence of the one-point statistics of the density field on the effective equation of state of the gas, the poor correlation of magnetic strength with density, the energy spectrum in weakly and highly compressible cases, and the one point statistics of the velocity field; 3) the effects of spectroscopic observation on distorting the physical structures and results from synthetic observations of the simulations, and 4) several dynamical and thermodynamical issues, such as the (apparently minor) role of the thermal instability in forming and confining clouds, the continuous, rather than abrupt, transition between ``phases'', which in turn may be consequences of the dynamics rather than the agents controlling it, the possibility of short time scales ($\\sim $ a few Myr) for molecular cloud formation, and the star-gas connection, mentioning that the models generally exhibit self-propagating star and cluster formation, while the stars may drive the medium- and small-scale gas motions, and that a ``star formation instability'' may induce chaotic behavior of the star formation rate locally. I conclude with a round-up view, and a discussion of the work needed ahead. ",
        "introduction": "Two of the most influential models of the interstellar medium (ISM) to date are the two- and three-phase models of Field, Goldsmith \\& Habing (1969) and of McKee \\& Ostriker (1977, hereafter, MO). These models relied on the then known atomic and radiative heating and cooling processes to provide a self-consistent picture of the ISM in which the concepts of thermal and pressure equilibria played a fundamental role. Another important model was the time-dependent model of the ISM by Gerola, Kafatos \\& McCray (1974), which was presented as an alternative to the pressure equilibrium two-phase model of Field et al.\\ (1969) but assumed radically different conditions: a constant-density medium under the influence of stochastic, local heating events that should cause strong local fluctuations of pressure and temperature, because the cooling and recombination times are comparable or shorter than the time between exposure of a given gas parcel to one of those heating events. Note that the MO model also recognized the existence of local fluctuations in the pressure, although it was still based on the premise of ``rough pressure balance''.  Nevertheless, both the equilibrium and the time dependent models left out a number of important aspects in the ISM physics. The multiphase equilibrium models essentially neglected the possibility of large pressure fluctuations in the ISM. The time dependent model instead included this possibility as a fundamental premise, but still neglected the fact that such fluctuations should induce motions, which should in general be turbulent and cause strong density fluctuations. The turbulence involves gas motions at all scales which not only provide ram ``pressure'', but also mixing, and can produce compressions rather than ``support'' (e.g., Elmegreen 1993; \\BP, \\VS\\ \\& Scalo 1999a).  Moreover, both the time-dependent and the MO models omitted other sources of pressure in the ISM, such as magnetic fields and cosmic rays. The pressure from these agents is in fact significantly larger than the thermal pressure (e.g., Boulares \\& Cox 1990). Elmegreen (1991, 1994) has performed a combined instability analysis including self-gravity, cooling and heating, and magnetic fields, but the full nonlinear behavior can only be dealt with by means of numerical simulations of the gas dynamics in the Galactic disk. In this review, I will summarize a variety of results in this area, pioneered by Bania \\& Lyon (1980), concerning the spatial distribution both on the Galactic plane and perpendicular to it (\\S 2), the one-point and correlation statistics of the physical fields (\\S 3), the comparison between ``synthetic'' observations of the simulations and actual observations of the ISM (\\S 4), and several dynamical aspects, including the role of thermal and ram pressures, the virial balance of clouds, the nonlinear Q parameter and the star-gas connection (\\S 5). I conclude (\\S 6) with an attempt to present a comprehensive scenario, which should constitute the first step towards a full dynamical theory of the ISM that can take over where the time-dependent and the MO models left off.  Given the focus on HI gas of the present conference, this review concentrates on hydro- and magnetohydrodynamic numerical simulations of the multi-temperature\\footnote{In this review the common term ``multi-phase'' will be avoided and substituted by ``multi-temperature'', since, as discussed in \\S 5.1, numerical simulations do not in general support the existence of sharp phase transitions in interstellar gas.} gas in the Galactic disk. Thus, the vast literature existing on numerical simulations of isothermal gas, aimed at the molecular-cloud regime, will necessarily be excluded. The reader is directed to the review by \\VS\\ et al.\\ (2000) for a discussion of that area current up to 1999, and that by Mac Low \\& Klessen (2002) for the more recent results. Non-hydrodynamical models of the dynamical systems kind have also been excluded (see, e.g., the review by Shore \\& Ferrini 1995). ",
        "conclusions": "In this review I have summarized a large body of results derived mainly from numerical simulations of the ISM, which have allowed workers to capture many of the complex dynamical aspects of ISM structure and evolution. The numerical models have suggested a much more complex ISM than a simple three-phase medium in pressure equilibrium. Instead, a turbulent continuum has emerged, apparently without sharp phase segregation, and in which turbulent ram pressure is mostly responsible for cloud formation within the large complexes, while strong thermal pressure imbalances due to local stellar energy injection are responsible for powering the turbulence through expanding shells, although otherwise the near constancy of the thermal pressure has little effect in confining density structures, since {\\it the absence of a pressure gradient does not imply that inertial motions cannot exist}. The medium is not simply a collection of overlapping shells with clouds forming in their compressed layers, but instead is globally turbulent, with shells and all structures in general ``morphing'' and  merging into the global turbulence. At the largest scales, combined gravitational, magnetic and thermal instabilities appear to contribute, together with spiral density waves and supershells, to the formation of the largest complexes and perhaps voids, as well as possibly feeding the turbulence from the largest scales.  The individual structures being transient, the appropriate description for the global structure is statistical. Some statistical measures of the flow (the density PDF and the energy spectrum) depend on the cooling ability of the flow, which in turn determines its compressibility. Since this in turn determines the ability to form gravitationally bound structures through compressions, and therefore stars which feed the turbulence, a full feedback loop clearly exists analogous to the old Oort (1954) cycle.  However, as this review probably makes evident, the body of results is highly scattered, and a coherent, global theory of the ISM is still lacking. Such a theory should be able to predict fundamental statistical indicators of the ISM such as its topological properties (i.e., the statistical distribution of the mass and other physical quantities), the rate of production of collapsing objects (the star formation rate and its efficiency) and the mixing rate of the processed chemical elements, as well as the thermal and radiative properties of the various temperature and density regimes, all as a function of simple input physics, such  as the total available mass and angular momentum, the atomic processes determining the cooling rates, and the energy injection per source. Note that, in principle, even ``parameters'' such as the energy injection rate and the scale of injection, should be derivable from the theory (or model), because the sources in this case are linked to the flow dynamics, as is the case of large-scale instabilities and of the star formation rate. Many of the physical variables are expected to be highly fluctuating locally, and the prediction of both average (see, e.g., Blitz, this volume) and typical fluctuation values is crucial. Simulations including all relevant physical ingredients (self-gravity, disk rotation, magnetic fields, cosmic rays, stellar energy input, spiral density waves, chemistry) and that can perform the necessary radiative transfer, all at high enough resolution, are needed. So are careful experiments which allow disentangling the effects of all those ingredients, and of the average quantities versus those of the fluctuations. Consideration of the gravitational effect of the stars is also needed, which implies that hybrid gas+stars similations are needed. Then, continuous feedback to/from observations is essential through the production of synthetic observations of the numerical models. A long, exciting way still lies ahead to a comprehensive theory that grasps the dynamics, thermodynamics and statistics of the ISM with the same level of detail with which the time-dependent and the multi-phase models grasped the thermal and radiative issues."
    },
    "0201/astro-ph0201302_arXiv.txt": {
        "abstract": "We report optical and IR observations of the ASCA X-ray pulsar system AX J0051-733. The relationship between the X-ray source and possible optical counterparts is discussed. Long term optical data from over 7 years are presented which reveal both a 1.4d modulation and an unusually rapid change in this possible binary period. Various models are discussed. ",
        "introduction": "High Mass X-ray binaries (HMXBs) are traditionally divided into Be/X-ray and supergiant binary systems. A survey of the literature reveals that of the 96 proposed massive X-ray binary pulsar systems, 67\\% of the identified systems fall within the Be/X-ray group of binaries.  The orbit of the Be star and the compact object, a neutron star, is generally wide and eccentric.  The optical star exhibits H$\\alpha$ line emission and continuum free-free emission (revealed as excess flux in the IR) from a disk of circumstellar gas. Most of the Be/X-ray sources are also very transient in the emission of X-rays.  The source that is the subject of this paper, AX J0051-733, lies in the Small Magellanic Cloud, a region of space that is extremely rich in HMXBs. It was reported as a 323s pulsar by Yokogawa \\& Koyama (1998) and Imanishi et al (1999). Subsequently Cook (1998) identified a 0.7d optically variable object within the ASCA X-ray error circle. The system was discussed in the context of it being a normal HMXB by Coe \\& Orosz (2000) who presented some early OGLE data on the object identified by Cook (1998) and modelled the system parameters. Coe \\& Orosz identified several problems with understanding this system, primarily that if it was a binary then its true period would be 1.4d and it would be an extremely compact system. In addition, the combination of the pulse period and such a binary period violates the Corbet relationship for such systems (Corbet, 1986).  In this paper we report on extensive new data sets from both OGLE and MACHO, as well as a detailed photometric study of the field. The results reveal many complex observational features that are hard to explain in the traditional Be/X-ray binary model.  \\section {X-ray source location}  As will be seen from the photometric results presented below, it is critical to establish the correct optical counterpart to the X-ray pulsar. In particular, it is vital to clearly link the ASCA source to an optical object, and other ROSAT X-ray sources may, or may not be relevant (because no pulsations have been detected from ROSAT objects). Figure~\\ref{optir} illustrates the somewhat complex situation associated with this object. In this figure the large dotted circle indicates the original ASCA X-ray position and uncertainty from Imanishi et al (1999). Within this error circle lie the much smaller error circles of the ROSAT sources RX J0050.8-7316 (Cowley et al, 1997) and RX J0050.7-7316 (Kahabka 1998). Subsequently, the position of the ASCA error circle was refined to the large solid circle shown in the figure (Imanishi 2001, private communication).    \\begin{figure} \\begin{center} \\psfig{file=fig1.eps,width=3in} \\end{center} \\caption{Finding charts for AX J0051-733 covering a field of 4 x 4 arcmin created using an optical V band image from this work. The northern ROSAT circle refers to RX J0050.8-7316 and the southern one is that of RX J0050.7-7316. \\label{optir}} \\end{figure}     Within the ROSAT error circle for RX J0050.7-7316 and the revised ASCA circle for AX J0051-733 lies an obvious optical object that has been proposed as the counterpart to both of these X-ray objects (Cowley et al, 1997, Schmidtke \\& Cowley, 1998, Coe \\& Orosz, 2000). It is a blue star exhibiting variability which strongly suggests that it is a Be star companion to the X-ray pulsar. This object also shows a strong 0.7d optical modulation (or possibly twice that value) which could be associated with a binary period of the system (Cook 1998, Coe \\& Orosz 2000). However, this period is very short for a HMXB and the modulation signature atypical of that seen from such objects.  Consequently, it was felt necessary to revisit the linking of this optical object with the ASCA pulsar to make sure that some other candidate was not more appropiate within the X-ray error circle. ",
        "conclusions": "Detailed optical observations and analysis have been carried out of the proposed counterpart to AX J0051-733. The most likely counterpart has been identified on the basis of its colours and H$\\alpha$ emission. However this object is revealed to have a strong 0.7/1.4d modulation from long-term MACHO and OGLE observations. Furthermore this strong period is shown to be changing at a rate of 13.5s/year. It is hard to reconcile all these observations with the classic Be/X-ray binary model and further studies of this system are urgently required."
    },
    "0201/astro-ph0201134_arXiv.txt": {
        "abstract": "We have developed a Monte Carlo technique to test models for the true power spectra of intermittently sampled lightcurves against the noisy, observed power spectra, and produce a reliable estimate of the goodness of fit of the given model.  We apply this technique to constrain the broadband power spectra of a sample of four Seyfert galaxies monitored by the {\\it Rossi X-ray Timing Explorer} ({\\it RXTE}) over three years.  We show that the power spectra of three of the AGN in our sample (MCG-6-30-15, NGC~5506 and NGC~3516) flatten significantly towards low frequencies, while the power spectrum of NGC~5548 shows no evidence of flattening.  We fit two models for the flattening, a `knee' model, analogous to the low-frequency break seen in the power spectra of BHXRBs in the low state (where the power-spectral slope flattens to $\\alpha=0$) and a `high-frequency break' model (where the power-spectral slope flattens to $\\alpha=1$), analogous to the high-frequency break seen in the high and low-state power spectra of the classic BHXRB Cyg~X-1. Both models provide good fits to the power spectra of all four AGN. For both models, the characteristic frequency for flattening is significantly higher in MCG-6-30-15 than in NGC~3516 (by factor $\\sim10$) although both sources have similar X-ray luminosities, suggesting that MCG-6-30-15 has a lower black hole mass and is accreting at a higher rate than NGC~3516. Assuming linear scaling of characteristic frequencies with black hole mass, the high accretion rate implied for MCG-6-30-15 favours the high-frequency break model for this source and further suggests that MCG-6-30-15 and possibly NGC~5506, may be analogues of Cyg~X-1 in the high state.  Comparison of our model fits with naive fits, where the model is fitted directly to the observed power spectra (with errors estimated from the data), shows that Monte Carlo fitting is essential for reliably constraining the broadband power spectra of AGN lightcurves obtained to date. ",
        "introduction": "\\label{intro} The strong and rapid X-ray variability observed in many Seyfert galaxies on time-scales of a day or less provides strong evidence that the X-rays are emitted close to the central black hole.  Early efforts to characterise the X-ray variability, using data from {\\it EXOSAT}, showed that it has a scale-invariant, red-noise form on time-scales from a few hundred seconds up to the few-days duration of the observations (M$^{\\rm c}$Hardy \\& Czerny 1987, Lawrence et al. 1987).  Later studies of the X-ray variability properties of large samples of radio-quiet AGN showed that the variability amplitude scales inversely with luminosity (Green, M$^{\\rm c}$Hardy \\& Lehto 1993; Lawrence \\& Papadakis 1993; Nandra et al. 1997). One possible explanation of this result is that the higher luminosity AGN contain more massive black holes and the variability time-scales in AGN scale with black hole mass. Intriguingly, black hole X-ray binary systems (BHXRBs) also show red-noise type variability of a similar amplitude to AGN, on time-scales less than seconds.  The similarity in X-ray variability properties of AGN and BHXRBs raises the possibility that the processes causing variability in AGN and BHXRBs are the same and that any characteristic variability time-scales scale with the central black hole mass. This possibility can be tested by comparing the detailed X-ray timing properties of BHXRBs and AGN. \\\\ Timing studies of BHXRBs are usually carried out in the frequency domain using the power spectrum, which shows the contribution of variations on different time-scales (corresponding to power-spectral frequencies) to the total variability of the lightcurve. The power spectra of BHXRBs are dominated by a broadband noise component (van der Klis 1995).  On short time-scales, the variability is characterised as scale-invariant `red noise', producing a power-law power spectrum (power $P(\\nu)$ at frequency $\\nu$ is given by $P(\\nu)\\propto \\nu^{-\\alpha}$ where $\\alpha$ is the power-spectral slope) of slope $\\alpha \\sim1$--$2.5$ (van der Klis 1995). In the `low' state, characterised by a relatively hard X-ray spectrum (similar to that of AGN), the power spectrum flattens towards lower frequencies so that on long time-scales the X-ray lightcurve becomes `white noise', with corresponding slope $\\alpha \\simeq0$. For example, in the classic BHXRB system Cyg~X-1, the power-spectral flattening is well described by a power-law with two breaks, a high-frequency break which varies between 1 and 6~Hz, above which the power-spectral slope varies between $\\alpha \\sim1.5$--$2.4$ and a low-frequency break which varies between 0.04 and 0.4~Hz, above which the power-spectral slope $\\alpha \\simeq1$ and below which the slope $\\alpha \\simeq0$ (Belloni \\& Hasinger 1990).  In contrast, the power spectrum of the `high' (soft energy spectrum) state seen in some BHXRBs (inluding Cyg~X-1) does not flatten to zero slope; instead, the slope $\\alpha\\simeq1$ below the high-frequency break extends to $<10^{-2}$~Hz (e.g. Cui et al. 1997). \\\\ In order to test the hypothesis that the X-ray variability of AGN is similar to that of BHXRBs over a broad range of time-scales, we must search for low-frequency flattening in the broadband power spectra of AGN. By fitting models with power-spectral breaks to the AGN power spectra and comparing the estimated break frequencies with what we expect if they correspond to similar breaks in BHXRB power spectra, we can test the possibility that the power-spectral shape is really the same and scales simply with black hole mass. If so, we expect break frequencies in AGN to be found at frequencies of $\\sim10^{-5}$~Hz or lower, so that monitoring observations on time-scales of weeks or longer are necessary to detect any flattening in the power spectrum. \\\\ Early attempts to measure broadband power spectra of AGN were hampered by the sparseness of long-term archival lightcurves, which had to be constructed from data obtained by several missions (M$^{\\rm c}$Hardy 1988). Nonetheless, some evidence for power-spectral flattening was found, but models for the form of the flattening could not be constrained (Papadakis \\& M$^{\\rm c}$Hardy 1995). \\\\ Ideally, broadband power spectra should be measured from lightcurves obtained with frequent and regular sampling over a long duration, which previous missions were not optimised to do.  The {\\it Rossi X-ray Timing Explorer} ({\\it RXTE}), which has just such a capability, was launched in December 1995.  {\\it RXTE} carries a large-area proportional counter array (the PCA) which can detect many AGN with good signal-to-noise in less than 1000~s, but most importantly, {\\it RXTE} can slew rapidly so that it may monitor many targets with frequent 1~ks snapshots. \\\\ We have monitored a sample of 4 Seyfert galaxies (MCG-6-30-15, NGC~4051, 5506 and 5548) with {\\it RXTE} since 1996, in order to measure their broadband power spectra.  These objects are known to be significantly X-ray variable and cover a broad range of X-ray luminosity (NGC~4051$\\sim5\\times 10^{41}$~erg~s$^{-1}$, MCG-6-30-15 and NGC~5506$\\sim1.5\\times 10^{43}$~erg~s$^{-1}$, NGC~5548$\\sim5\\times10^{43}$~erg~s$^{-1}$) and presumably, a broad range of black hole masses.  We describe the power spectrum of NGC~4051, which shows unusual non-stationarity in its lightcurve (Uttley et al. 1999) in a separate paper (Papadakis, M$^{\\rm c}$Hardy \\& Uttley, in prep.).  In this paper, we present a power-spectral study of the remaining three objects in our sample using data from {\\it RXTE} cycles 1, 2 and 3, also including the excellent lightcurves obtained as part of a separate power-spectral study of the Seyfert~1 galaxy NGC~3516 (luminosity$\\sim1.5\\times 10^{43}$~erg~s$^{-1}$), by Edelson \\& Nandra (1999). We describe data reduction and present the lightcurves in Section~2. \\\\ The estimation of the underlying power-spectral shape from lightcurves which are discretely (and possibly unevenly) sampled is hampered by the distorting effects of aliasing and red-noise leak.  A further serious problem is that the measured power spectra are intrinsically noisy, and reliable errors on the power in each frequency bin cannot be estimated from the data (especially at low frequencies), due to the small number ($<20$) of power-spectral measurements made in each frequency bin.  In Section~3, after presenting the observed power spectra, we describe these problems, which previous efforts to constrain the shape of the broadband power spectrum of AGN using {\\it RXTE} data (e.g. Edelson \\& Nandra 1999, Nowak \\& Chiang 2000) have not accounted for. \\\\ To overcome the difficulties in estimating the true power-spectral shape, we have developed a method which we call {\\sc psresp}, based on the response method (Done et al. 1992) which uses Monte Carlo simulations of lightcurves to take account of the distorting effects of sampling and to estimate uncertainties, allowing us to test various power-spectral models against the data.  We describe {\\sc psresp} in Section~4, and apply it to the lightcurves of our sample of Seyfert galaxies in Section~5, in order to test for flattening in their broadband power spectra and constrain simple models for describing any flattening we see.  In Section~6, we compare our results with those obtained by naively fitting the observed power spectrum (without taking proper account of errors and the distortion due to sampling), use our power spectral measurements to estimate the black hole masses of the AGN in our sample and discuss some of the implications of our results, before making concluding remarks in Section~7. ",
        "conclusions": "We have presented long-term {\\it RXTE} monitoring lightcurves for 4 Seyfert galaxies, MCG-6-30-15, NGC~5506, NGC~5548 and NGC~3516, and measured their broadband power spectra to determine if they flatten towards low frequencies, like those of BHXRBs.  The interpretation of power spectra measured from discretely (and sometimes unevenly) sampled lightcurves is complicated by the distorting effects of red-noise leak and aliasing.  A further serious problem is that reliable errors in the power in each frequency bin cannot be defined from the data, due to the small number ($<20$) of power-spectral measurements in each frequency bin.  \\\\ To overcome these difficulties, we have built on the response method of Done et al. (1992), to develop a reliable Monte Carlo method for testing models for the underlying power-spectral shape of discretely, unevenly sampled lightcurves.  Our method, {\\sc psresp}, takes proper account of the effects of aliasing and red-noise leak and crucially, uses the distribution of simulated power spectra to define reliable confidence limits on our model fits. \\\\ We have used {\\sc psresp} to test simple models for the power-spectral shape of the active galaxies in our sample.  Our main results are: \\begin{enumerate} \\item[1.] A single power-law model for the power spectra, with no low-frequency flattening, is rejected at better than 95\\% confidence for MCG-6-30-15, NGC~5506 and NGC~3516.  The power spectrum of NGC~5548 is consistent with a single power-law. \\item[2.] Both the knee model (flattening to slope $\\alpha=0$) and high-frequency break model (flattening to $\\alpha=1$) provide good fits to the power spectra of all four sources.  Knee/break frequencies are well constrained for all sources except NGC~5548, for which we can define upper limits only. \\item[3.] The characteristic knee/break frequency measured for MCG-6-30-15 is significantly higher (99\\% confidence) than the corresponding frequency in NGC~3516, even though both sources have similar X-ray luminosities ($\\sim1.5\\times 10^{43}$~erg~s$^{-1}$). \\item[4.] If the knees or breaks correspond to those seen in the low-state power spectrum of Cyg~X-1, and the characteristic knee/break time-scales scale linearly with black hole mass, the black hole masses estimated from the knee model are lower than those estimated by the break model, but remain consistent with sub-Eddington accretion rates.  In the case of MCG-6-30-15, the conservative lower limit on accretion rate estimated from the knee model is an order of magnitude higher than that seen in the low state of Cyg~X-1, favouring the high-frequency break model and further suggesting that the break seen in the power spectrum of MCG-6-30-15 (and possibly NGC~5506) is analogous to the high-frequency break seen in the {\\it high state} power spectrum of Cyg~X-1. \\item[5.] The $\\nu P(\\nu)$ power spectra of the Seyfert galaxies studied here are similar in amplitude to the $\\nu P(\\nu)$ power spectra of Cyg~X-1 in the high and low states. \\end{enumerate} Conclusions 4 \\& 5 imply that the power spectra of AGN are consistent with being identical in shape and fractional RMS amplitude (integrated over the whole power spectrum) to those of BHXRBs, with characteristic time-scales scaling linearly with black hole mass. Arguments based on accretion rate seem to favour the high-frequency break model over the knee model in at least one source (MCG-6-30-15), although this evidence is circumstantial.  Further monitoring observations (which are currently underway) are needed to reject either model on the basis of power-spectral measurements alone, and to confirm the interesting possibility that two of the objects in our sample are analogous to high-state BHXRBs.  \\subsection*"
    },
    "0201/astro-ph0201378_arXiv.txt": {
        "abstract": "We have carried out a survey for `giant pulses' in 6 young, Vela-like pulsars. In no cases did we find single pulses with flux densities more than 10 times the mean flux density. However, in PSR B1706--44 we have detected giant micro-pulses very similar to those seen in the Vela pulsar. In PSR B1706--44 these giant micro-pulses appear on the trailing edge of the profile and have an intrinsic width of $\\sim$1 ms. The cumulative probability distribution of their intensities is best described by a power-law. If the power-law continues to higher intensities, then $3.7\\times 10^6$ rotations are required to obtain a pulse with 20$\\times$ the mean pulse flux. This number is similar to the giant pulse rate in PSR B1937+21 and PSR B1821--24 but significantly higher than that for the Crab. ",
        "introduction": "Despite three decades of intensive study, the mechanism producing pulsar radio emission is poorly understood. Fluctuations in the intensity of the radio radiation provide important constraints on plausible mechanisms. Single-pulse studies of bright pulsars detect a variety of patterns in the intrinsic intensity fluctuations, including nulling and drifting phenomena.  The distribution of integrated pulse energies, however, has only a modest dispersion. Johnston et al. (2001)\\nocite{jvkb01} showed that in the Vela pulsar, 99.5\\% of all pulses are within a factor of 3 of the mean flux density, $\\langle S \\rangle$ and that the histogram of pulse energies is a Gaussian when plotted in the log. This distribution seems typical for most pulsars \\cite{hw74,rit76}.  In contrast, the Crab pulsar emits pulses with flux densities $> 20 \\times \\langle S \\rangle$, extending up to $> 2 \\times 10^3 \\langle S \\rangle$ \\cite{lcu+95} which were instrumental in the original detection of the Crab \\cite{sr68}. These giant pulses are typically broadband \\cite{mof97,sbh+99} and of short duration, with widths of order a few $\\mu$s and structure down to 10~ns \\cite{han96b}. They are localized to the main and interpulse phase windows and follow an intensity distribution best characterized as a power law with index $\\sim 3-3.5$.  The discovery of similar pulses from the millisecond pulsar PSR B1937+21 \\cite{sb95,cstt96} was surprising. The pulses are extremely short ($\\tau < 0.3 \\mu$s) events confined to small phase windows trailing the main pulse and interpulse, again with an approximately power-law distribution of pulse energies \\cite{kt00}. Since PSR B1937+21 is the only known radio pulsar with an estimated magnetic field at the light cylinder larger than that of the Crab, it has been suggested that this is a key parameter controlling the giant pulse phenomenon \\cite{cstt96}.  Johnston et al. (2001)\\nocite{jvkb01} have recently found that a small subset of pulse phases from the Vela pulsar have a very wide distribution of peak fluxes. The pulses are localized to a phase window $\\sim$1 ms prior to the bulk of the integrated pulse emission, are of short duration and are highly polarized. Johnston et al. (2001) called these giant micro-pulses. It is not clear if these events are related to true giant pulses; the largest Vela pulses observed to date have $ S < 10\\langle S \\rangle$, but these narrow pulses have peak fluxes exceeding $40\\times$ the integrated peak intensity. Kramer, Johnston \\& Van Straten (2001)\\nocite{kjv01} have shown that these giant micro-pulses have a power-law distribution and the extended tail of the distribution may continue into the true giant pulse regime. Cairns, Johnston \\& Das (2001)\\nocite{cjd01}, in contrast, recently showed that a log-normal distribution provided an excellent fit to the flux densities in individual phase bins across the main peak of the Vela profile. It seems likely that the same distribution is also applicable to other pulsars. Therefore a potential discriminator of giant pulse activity is the change from a log-normal distribution to a power-law one.  To explore the connection between the giant pulses in the Crab and PSR B1937+21 and the large individual pulses in the Vela pulsar, we obtained fast time-sampled data for several young and millisecond pulsars. In an earlier paper \\cite{rj01} we reported the detection of giant pulses from the millisecond pulsar with the next highest known light cylinder field, PSR B1821$-$24. In this paper we report on our survey of young pulsars focussing on the results obtained for PSRs B1046--58 and B1706--44. ",
        "conclusions": "In the Vela pulsar, log-normal statistics are adequate to fit the distribution of flux across the bulk of the pulse profile (Kramer et al. 2001, Cairns et al. 2001)\\nocite{kjv01,cjd01} but the width of the distribution is significantly larger at the edge of the profile than in the middle. In addition to this, there are giant micro-pulses which occur well before the main pulse phase. These giant micro-pulses have a half-width of $\\sim$200 $\\mu$s, and are not at a fixed phase, but have an inherent `jitter' of about 1 ms \\cite{jvkb01}. Their distribution is best described by a power-law (Kramer et al. 2001)\\nocite{kjv01}. We have found new examples of both these phenomena in our current survey.  PSR B1706--44 shows an additional example of giant micro-pulses. This time, however, the giant micro-pulses are located on the trailing edge of the pulse, and are somewhat wider then in Vela with a half-width close to 1 ms. Again there is some phase jitter as to the location of the pulse maximum. For this pulsar also, the distribution of fluxes is clearly power-law at high amplitudes.  In PSR B1046--58 we clearly see large amplitude pulses on both the leading and trailing edges of the integrated pulse profile. The intensity distribution at these phases, however, is adequately described by a log-normal distribution with moderate width. The peak fluxes achieved are in excess of 20$\\times$ the mean flux density but there is no evidence for a power-law tail to the distribution. In this regard they are similar to the fluctuations seen on the rising edge of the Vela pulsar.  No giant micro-pulses are detectable in PSRs J1105--6107, J1420--6048, B1509--58 or J1617--5055. However, the sensitivity to giant micro-pulses in these pulsars is not as good as for PSR B1706--44 or B1046--58. To demonstrate this, let's assume that any giant pulses would be similar to those seen in PSR B1706--44; i.e. they would have an intrinsic width of 0.5 ms, a power-law index of --3.0 and occur every $10^7$ rotations. The brightest giant pulse in the observing time for PSRs J1105--6107, J1420--6048, B1509--58 and J1617--5055 would therefore be less than 500 mJy (in 0.5 ms) and not detectable in the noise. Giant pulses could only be detected in these pulsars if they were intrinsically very narrow, and/or occured much more frequently than in either Vela or PSR B1706--44.  We find, as in previous studies (e.g. Taylor et al. 1975\\nocite{tmh75}), that the modulation of pulse intensity is larger in the wings of the profile of Vela-like pulsars. PSR B1046--58 is a good example of this behaviour: its phase resolved intensity distributions are well described by a log-normal distribution whose width increases towards the profile edge. The giant micro-pulses in Vela and now PSR B1706--44 appear to represent a distinct pulse population, with emission phases well separated from the bulk of the integrated pulse profile. Both of these phenomena represent enhanced intensity fluctuations.  Observations of the Crab pulsar at infra-red, optical and higher energies (e.g. Lundgren et al. 1995, Patt et al. 1999, Romani et al. 2001) \\nocite{lcu+95,puz+99,rmc+01} show that there are no detectable intensity fluctuations in the (incoherent) high energy emission associated with the radio giant pulses. Thus we must conclude that giant pulses represent a variation in the coherence of the particle distribution. An important question is why coherence fluctuations appear to be strongest at the pulse edges.  One clue may be found in the connection of pulse width to altitude. Crudely, in a dipole polar cap of radius $r$, the width of the edge of the open zone maps as $\\delta \\phi \\propto r^{1/2}$. If interpreted as the edge of a polar zone, the leading and trailing edges of the profiles where the modulation index increases represent locations 1.5-2$\\times$ higher than that of the mean pulsar emission. If we adopt the classical polar gap picture, with a pair formation front producing a dense plasma at a few stellar radii, then we can identify pulse intensity fluctuations with the growth of instabilities in the outflowing pair wind. Larger widths represent higher altitudes, and in this picture larger phase-space amplitudes for the instability-driven fluctuations, which have had more time to grow in the outflow. Alternative pictures exist, such as the slot-gap scenario of Arons (1983)\\nocite{aro83b} for which the gap height is larger at the edge of the polar cap. Again there could be a plausible connection with increased instability.  In both PSR B1706--44 and Vela the giant micro-pulses are located far from the peak of the pulse profile, representing an emission altitude $\\sim$4$\\times$ higher. It is plausible that instability growth reaches a different regime at these heights. Clearly a true power-law intensity distribution suggests that the instabilities have left the (linear) stochastic growth regime (Cairns et al. 2001). Power-law behaviour suggests a causal connection between different intensity scales such as a cascade from one energy scale to another, or a self-similar coupling as in a saturated growth scheme. Such processes might develop late (higher) in the polar wind outflow.  The giant micro-pulse emission is however well separated from the bulk of the radio emission in these objects, suggesting instead an independent origin. There is good evidence that for the Crab and possibly PSR B1821--24, the giant pulse radio emission is co-located with high energy emission in the outer magnetosphere, where strong power law X-ray components indicate sites of dense pair production. Vela and PSR B1706-44 are also $\\gamma$-ray emitters; in the outer magnetosphere picture this high energy emission comes from the pole opposite to that viewed in the radio. They do not show strong X-ray pulses and so it is perhaps not surprising that we have not found giant pulse emission coincident with the $\\gamma$-rays in these objects, as the Earth line-of-sight evidently does not sample regions of dense pair production. Nevertheless, there may be a high energy connection to the giant micro-pulse emission if the outer magnetosphere above the radio pole is also actively producing pairs. The high energy emission from beyond the null charge surface would not be visible from this pole, but some pairs from the gap would be expected to flow inward past the null charge surface. We can speculate that this plasma mirrors in the converging field lines above the radio pole and that this mirrored, counter-streaming population would suffer instability growth and produce the giant micro-pulse components.  A final question concerns the asymmetry of the giant micro-pulse component. Why does it lead in the case of Vela, but lag for PSR B1706$-$44? For Vela, the radio pulse has a steep rise and slow fall suggesting the leading edge of a cone; the giant micro-pulse component is at higher altitudes on the same side. For PSR B1706$-$44, one would then infer that the main radio pulse is the trailing edge of a cone, although this is morphologically less clear. If there is a high energy connection, then it is intruiging to note that the leading $\\gamma$-ray pulse is stronger for Vela, while for PSR B1706$-$44, the trailing component is stronger at most high energies. More examples are clearly needed to see if the dominance of leading versus trailing giant micro-pulses has a deterministic connection with other pulsar emission and if some global asymmetry in the magnetosphere geometry controls this choice."
    },
    "0201/astro-ph0201187_arXiv.txt": {
        "abstract": "We model in simple terms the angular-momentum problems of galaxy formation in CDM cosmologies, and identify the key elements of a scenario that may solve them.  The buildup of angular momentum is modeled via dynamical friction and tidal stripping in a sequence of mergers. We demonstrate how over-cooling in incoming halos leads to a transfer of angular momentum from the baryons to the dark matter, in conflict with observations. By incorporating a simple recipe of supernova feedback, we are able to solve the problems of angular momentum in disk formation.  Gas removal from the numerous small incoming halos, which merge to become the low specific angular momentum ($j$) component of the product, eliminates the low-$j$ baryons.  Heating and puffing-up of the gas in larger incoming halos, combined with efficient tidal stripping, reduces the angular momentum loss of baryons due to dynamical friction. Dependence of the feedback effects on the progenitor halo mass implies that the spin of baryons is typically higher for lower mass halos.  The observed low baryonic fraction in dwarf galaxies is used to calibrate the characteristic velocity associated with supernova feedback, yielding $\\vfb \\sim 100\\kms$, within the range of theoretical expectations.  We then find that the model naturally produces the observed distribution of the spin parameter among dwarf and bright disk galaxies, as well as the $j$ profile inside these galaxies.  This suggests that the model indeed captures the main features of a full scenario for resolving the spin crisis. ",
        "introduction": "\\label{sec:intro}  The `standard' model of cosmology, which assumes hierarchical buildup of structure in a universe where the mass is dominated by cold dark matter (CDM), seems to be facing intriguing difficulties in explaining some of the robust observed properties of galaxies. Standing out among these problems is the inability of galaxy formation models to reproduce both the sizes and structure of disk galaxies. In hydrodynamical simulations baryons lose a significant fraction of their angular momentum leading to overly small disks. To avoid this, analytic and semi-analytic models commonly assume there is no angular momentum loss.  Recent studies, however, show that this leads to a variety of discrepancies with observations. Taken together this suggests a crisis in our understanding of the role of angular momentum in forming disk galaxies. Our aim here is to make progress in the effort to resolve this crisis by first reproducing the problems using a simple model in which the important physical ingredients are spelled out and well understood. Subsequently we incorporate an additional process into this model which then reproduces the observed sizes and structure of galactic disks.  The sizes of galactic disks are commonly linked to the angular momentum of their parent dark-matter halos \\citep{fe:80}. This modeling is based on the distribution of halo spin parameters as found in N-body simulations \\citep[see][and references therein]{bull:01b}, combined with the assumptions that the baryons and dark matter initially share the same distribution of specific angular momentum, $j$, within the halos \\citep[as seen in simulations by][]{bosch:02} and that $j$ is conserved as the baryons contract to form the disk \\citep[as suggested by][]{mest:63}. The sizes of disks obtained under these assumptions are roughly comparable to those observed.  However, high-resolution simulations that incorporate gas processes find this scenario to be invalid.  In particular, they find that a significant fraction of the angular momentum of the baryons is transfered to the dark matter, resulting in disk sizes roughly an order of magnitude smaller then those observed \\citep{ns:00,slgv:99,ns:97,nfw:95,nb:91}. This has been refered to as the {\\it angular momentum catastrophe}.  The angular momentum catastrophe is commonly associated with the problem of ``over-cooling\" in CDM-type scenarios \\citep{wf:91,wr:78}. Without sufficient feedback much of the gas cools quickly, contracts into small halos and then spirals deep into the centers of bigger halos, efficiently transferring its orbital angular momentum to the dark matter \\citep{ns:00}. It has therefore been speculated that enough energy feedback from supernova may prevent this over-cooling \\citep{lars:74,wr:78} and thus reduce the angular-momentum loss. Indeed, simulations where gas cooling is artificially suppressed till $z=1$ do not suffer from the angular momentum catastrophe \\citep{eew:00,wee:98}. Furthermore, simulations including some forms of feedback show a reduction in the amount of angular momentum lost \\citep{slgv:99}. However, while feedback has been studied using simplified approximations \\citep[e.g.,][hereafter DS]{ds:86}, a realistic implementation of feedback has proved challenging \\citep[see][and references therein]{tc:00}, though some partial progress may have been made recently \\citep{ft:00,tc:01,sh:02}. At this point, the feedback scenario has not yet been studied in satisfactory detail, nor has it been confirmed to solve the spin problem, or properly understood in basic terms. This motivates an attempt to understand the scenario and how it may work using a simple semi-analytic model, in which the basic elements are easily understood.  Even if the angular momentum catastrophe can be avoided such that the total baryonic spin agrees with the total dark-halo spin, there remain discrepancies between the angular-momentum properties of dark halos in N-body simulations and those of observed disk galaxies. Foremost among these is the apparent mismatch of the distribution of specific angular momentum within galaxies (loosely termed ``the $j$ profile\") between observations and N-body simulations. \\citet[hereafter BD]{bull:01b}  revealed a universal profile with an excess of both low-$j$ and high-$j$ material compared to the observed $j$ distribution in exponential galactic disks.  Subsequently, \\citet*[hereafter BBS]{bbs:01} have measured the $j$ distribution in a sample of dwarf disk galaxies, confirming in detail, case by case, the differences between the $j$ profile predicted by the simulated halos and those found in disk galaxies. We refer to this second discrepancy as the {\\it mismatch of angular-momentum profiles}.  Another spin problem is indicated by the observations of \\citet{dl:00}, who found that the scatter in the log of the spin parameter inferred from fitting the spread of observed disks sizes in a large sample of late-type disk galaxies is only $0.36\\pm0.03$, while the scatter seen in simulated halos is significantly larger, $0.5\\pm0.05$ (BD and references therein). This could have been explained by a scenario in which the disks formed in low-spin halos are unstable and thus become early-type galaxies, which implies that disk galaxies occupy only the higher spin halos \\citep*{mmw:98,bosch:98}. However, this scenario is in apparent conflict with the finding in N-body simulations that recent major mergers, usually identified with large spheroidal components, actually give rise to halos with higher spin than average \\citep{gard:01,wech:01b}.  This demonstrates that the exact connection between the angular-momentum properties observed in galaxies and in N-body simulations is unclear.  We propose that the solution to the crisis is spin segregation --- that the angular momentum distribution of baryons differs from that of the dark matter due to gas processes. These processes can either decrease or increase the specific angular momentum of the baryons relative to the dark matter.  Gas cooling generally results in lower spin for the baryons compared to the dark matter, but heating due to feedback reduces this effect, and gas removal in small halos actually results in a higher spin for the baryons compared to the dark matter.  In this paper we work out a simple model to explore these effects. Knowing that in a hierarchical scenario the halo buildup can be largely interpreted as a sequence of mergers between smaller halos, our model is based on a simple algorithm for the buildup of halo spin by summation of the orbital angular momenta of merging satellites \\citep{mds:02,vitv:01}. This algorithm has been found to match well the spin distribution among halos in N-body simulations.  In this paper we extend the model and find that it also reproduces the angular-momentum profile within halos. It therefore provides a useful insight into the origin of the spin problems, and a clue for their possible solution. We therefore use this model as a tool for incorporating the relevant baryonic processes, and especially feedback.  In \\se{halos} we describe the model for angular-momentum buildup in halos by mergers.  In \\se{oc} we model the angular momentum catastrophe as resulting from over-cooling and dynamical friction. In \\se{fb} we introduce a simple model of feedback. In \\se{obs} we study the resultant angular-momentum properties of the baryonic component both in bright and dwarf disk galaxies and compare them to observations. In \\se{models} we test the robustness of our results to the details of the feedback model assumed. We conclude and discuss our results in \\se{conc}. ",
        "conclusions": "\\label{sec:conc}  We used a simple model to address the two issues that create the angular-momentum crisis of galaxy formation within the CDM scenario. One, the angular momentum catastrophe is that the spin of the baryonic component in galaxies is typically comparable to that of the dark-matter halo (in bright galaxies) or even larger (in dwarf galaxies), while simple theoretical arguments and simulations involving gas dynamics predict a significant spin loss by the baryons due to over-cooling in merging halos. The other, the mismatch of angular-momentum profiles, is that the baryons in each galaxy tend to lack the low-$j$ tail (and the high-$j$ tail) of the distribution as predicted by simulations for the dark halos.  In an earlier paper \\citep{mds:02}, we showed that a simple algorithm, based on adding up the orbital angular momenta of the mergers in random realizations of merger histories, can successfully reproduce the distribution of spins among dark-matter halos as measured in N-body simulations of the $\\lambda$CDM cosmology. We showed here that an extension of such a model also reproduces the characteristic angular-momentum profile, i.e., the distribution of specific angular momentum within halos. This provided the basic tool for addressing a possible resolution to the spin problems by incorporating the effects of supernova feedback on the gas in halos before they merge into bigger halos --- a process we have termed spin segregation.  A simple analysis of how the orbital angular momentum in a merger turns into a spin profile suggested how feedback effects in the satellite before the merger event can eliminate the problem of spin loss.  The effective size of the gas component within the incoming satellite determines its tidal stripping position in the halo and thus the final spin that it will be left with after the merger. The finding, using the orbital-merger model, that the low-end tail of the $j$ distribution originates in many minor mergers, that tend to cancel each other's angular momentum, provided a possible solution to the spin-profile mismatch problem. The blowout of gas from small incoming halos would eliminate the low-$j$ tail of the baryon distribution in the merger product, as observed. The blowout has a particular strong effect in dwarf galaxies, because they are made of smaller progenitors that tend to loose more of their gas.  This results in a higher spin parameter for the baryons than the dark matter, as observed.  We constructed a simple semi-analytic model for simulating this process using a simplified model for the effects of feedback as a function of halo mass, including heating and blowout. For a given choice of exponent $\\gamma_2$, the model has one free parameter, the characteristic halo virial velocity $\\vfb$ for which the energy inputted to the IGM by supernova is sufficient to counter the effect of cooling. By matching the low baryonic fraction in the dwarf galaxies observed by BBS, we found that for $0.8 \\leq \\gamma_2 \\leq 3$ the characteristic velocity has to be in the range $60 \\leq \\vfb \\leq 130 \\kms$, which falls within the range of theoretical predictions (e.g., DS). We then found considerable agreement between the model predictions and the observed data for both the distribution of the spin parameter and the angular-momentum profile of baryons in dwarf and bright galaxies. We also noted that the same basic model may explain the other unresolved issues concerning angular momentum in galaxies, such as the spread in observed disk sizes and the identification of halos that form late type galaxies. The success of the model in matching several independent observations indicates that this simple model indeed captures the main features necessary for a full scenario involving feedback and mergers that can resolve the spin crisis in more detail. The next natural step should be to incorporate a more sophisticated feedback recipe into the model using the full machinery of semi-analytic models of galaxy formation. This will be a step towards the long-term goal of implementing feedback in full-scale hydrodynamical cosmological simulations.  Our work is an attempt to resolve the angular-momentum problems within the standard framework of CDM cosmology, which is so successful on large scales, using the effects of feedback which we know must occur. Another approach is to appeal to a different cosmological scenario, WDM, where the dark-matter particles are ``warm\" rather than ``cold\" \\citep{hogan:99,hd:00,pp:82}. WDM is less robust than CDM because it requires fine-tuning of a new parameter, the particle mass, which is constrained to be $\\simeq 1~keV$. Still, it is worth investigating since it may remedy the angular-momentum problems without appealing to strong feedback effects. The main distinguishing feature of the WDM scenario is that the formation of small halos is significantly suppressed, such that the validity of the explicit picture of halo buildup by the hierarchical build up of small halos becomes limited. Despite this difference from CDM, an N-body simulation of WDM \\citep{bkc:01} indicates that the angular-momentum properties of halos remain basically unchanged. This is not very surprising because in \\citet{mds:02} we found that the same angular-momentum properties can also be interpreted as a general result of tidal-torque theory, independent of the explicit picture of mergers (see also BD). In the absence of small halos at early times, one may expect less over-cooling in halos before they merge into other halos and thus less angular-momentum loss by the baryons. Indeed, hydrodynamical simulations of this scenario \\citep{sld:01} indicate that the angular momentum catastrophe is significantly remedied. However, the angular-momentum profile mismatch is still expected to be valid in WDM, and, in the absence of small halos, the feedback effects are weaker and may not be enough for resolving the problem. Thus the $j$-profile mismatch may be a crucial discriminator between solutions to the problems of CDM.  It is clear that some sort of heating may provide the cure for another problem of CDM, the dwarf satellite problem \\citep{klyp:99b,moore:99a}, where the predicted number of dwarf halos is much larger than the observed number of dwarf galaxies. \\citet{bkw:00} demonstrated that cosmological photo-ionization due to feedback from UV sources such as early stars and quasars can solve this problem. Feedback from supernova would qualitatively have a similar effect. While the number of dwarf satellites is automatically suppressed in WDM, it seems that the inclusion of minimum realistic feedback effects would reduce the predicted number of galaxies to significantly below the observed number, and thus be an overkill (J. Bullock, private communication).  Furthermore, it is becoming clear (Dekel et al., in preparation) that the key elements of our model, namely the tidal effects in mergers and the feedback effects in small halos, are also very relevant in understanding and resolving the third problem of CDM. This is the cusp/core problem, where the halos in simulations typically show steep cusps in their inner profiles \\citep{nfw:95,moore:99a}, while observations indicate flat cores at least in some low-surface-brightness galaxies \\citep{bmr:01}. An analysis of tidal effects explains the necessary formation of an asymptotic cusp in halos as long as satellites continue penetrating into the halo center. Feedback effects may puff up small satellites and prevent them from penetrating the core, thus allowing a stable core. These studies together indicate that our model indeed grasps the relevant elements of the complex processes involved, and that feedback effects may indeed provide the cure to all three major problems of galaxy formation in CDM.   This research has been supported by the Israel Science Foundation grant 546/98, by the US-Israel Binational Science Foundation grant 98-00217, and by the German-Israeli Science Foundation grant I-629-62.14/1999. We thank James Bullock and Frank van den Bosch for stimulating discussions. We also thank Frank van den Bosch for providing us with the BBS data in \\Fig{amd}."
    },
    "0201/astro-ph0201464_arXiv.txt": {
        "abstract": "The sensitivity of stellar spectra to $\\alpha$/Fe abundance changes is investigated with the aim to be detected photometricly and employed for scientific goals of the GAIA mission. A grid of plane parallel, line blanketed, flux constant, LTE model atmospheres with different [$\\alpha$/Fe] ratios was calculated. As a first step, the modelled stellar energy fluxes for solar-type stars and giants were computed and intercompared. The spectral sensitivity to $\\alpha$/Fe abundance changes is noticeable and has to be taken into account when selecting photometric filters for GAIA. The Ca\\,{\\sc ii} H and K lines and Mg\\,{\\sc i}\\,b triplet are the most sensitive direct indicators  of $\\alpha$/Fe abundance changes. ",
        "introduction": "The evolution of a galaxy is closely related with a gradual chemical enrichment. The enrichment and spatial distribution of the chemical elements depend on various galactic and stellar processes. In particular, the star formation (SF) history, the time delay between SF and the enrichment of the interstellar medium (ISM), the metal dependency of the nucleosynthesis, the galactic gas flows, and the mixing processes in the ISM are important. Since the individual elements are produced at various sites and on different time scales, the observed abundances are very useful in describing galactic evolution. The main archeological tracers of the chemical evolution are the elements produced on a short time scale ($10^7$ years) by massive stars ending as core-collapse supernovae (here called SN~II) and on a longer time scale ($10^9$ years) by Type~Ia (SN~Ia) supernova events. SNe~II contribute to the enrichment of the interstellar medium mainly with elements produced by the capture of $\\alpha$-particles ($\\alpha$-elements) and from the $r$-process, and SNe~Ia predominantly produce elements belonging to the Fe peak. Consequently, one of basic tools to constrain the evolution of a galaxy is the analysis of relations between ratios of [$\\alpha$-element/Fe] and Fe abundances [Fe/H]\\footnote {In this paper we use the customary spectroscopic notation [X/Y]$\\equiv \\log_{10}(N_{\\rm X}/N_{\\rm Y})_{\\rm star} - \\log_{10}(N_{\\rm X}/N_{\\rm Y})_\\odot$} for stars born at different times and in different parts of a galaxy. E.g., the theoretical evolutionary model of the Milky Way galaxy recently proposed by Chiappini et al.\\ (2001) predicts a slight decrease with distance in the average [$\\alpha$/Fe] ratios in stars born in the Galactocentric distance range 4--10 kpc and an increase with distance of this ratio in the range 10--18 kpc.  A first glance at the temporal behavior of $\\alpha$-elements shows that most of the metal-poor stars in the Galaxy appear to have been formed with enhanced abundances of oxygen and other $\\alpha$-elements (i.e., Ne, Mg, Si, S, Ar, Ca and Ti). For stars with [Fe/H]$\\le-1$, the mean value of [$\\alpha$/Fe] lies between $+0.3$ and $+0.4$, with no discernible dependence on metallicity (cf. Pagel and Tautvaisiene 1995; Samland 1998). A more precise analysis shows that there is a significant population of field stars with [$\\alpha$/Fe]$\\approx 0.0$ (see Fig.4 by Nissen and Shuster, 1997). Surveys by Shuster et al.\\ (1993) and Carney et al.\\ (1996) report evidence for about $0.1-0.2$\\,dex variations in the [$\\alpha$/Fe] value at a fixed [Fe/H]. Carney et al.\\ (1997) have found that the high-velocity subgiant with apogalacticon distance over 20 kpc BD\\,$+80^{\\circ} 245$ has $<[\\alpha/{\\rm Fe}]>=-0.29\\pm 0.02$ despite to its low metallicity, [Fe/H]=$-1.86$. At the same time there are metal deficient giants with [$\\alpha$/Fe] of +0.7\\,dex as reported by Giridhar et al.\\ (2001). It is important to map the abundance pattern of $\\alpha$-elements in the Galaxy and understand the origin of variations.  A central element of the GAIA mission is the determination of the star formation histories, as described by the temporal evolution of the star formation rate, and the cumulative numbers of stars formed in the bulge, inner disk, Solar neighbourhood, outer disk and halo of our Galaxy (Perryman et al.,\\ 2001). Such information, together with the kinematic information from GAIA, and complementary chemical abundance information, again primarily from GAIA, may give us the full evolutionary history of the Galaxy. Knowledge of information on $\\alpha$/Fe abundance ratious in stars is very important for their age determination (cf. VandenBerg et al., 2000; Salasnich et al., 2000). In our study, a first attempt is made to investigate the sensitivity of stellar spectra to $\\alpha$/Fe abundance variations and their detection by GAIA photometry. ",
        "conclusions": "The spectral flux sensitivity to $\\alpha$/Fe abundance changes is noticeable and has to be taken into account in GAIA photometry. The Ca\\,{\\sc ii} H and K lines and Mg\\,{\\sc i}\\,b triplet are most sensitive direct indicators of $\\alpha$/Fe abundance changes and might be used for the photometric determination of $\\alpha$-element abundances. The photometric systems proposed for GAIA have to be carefully tested for accounting of the effects of $\\alpha$-element abundance variations and their determination.  Photometric classification of stars should provide as many physical parameters as possible. Depending on the accuracy with which the fundamental parameters are known, we should seek to determine abundances not only of $\\alpha$-elements but of carbon and nitrogen as well."
    },
    "0201/astro-ph0201008_arXiv.txt": {
        "abstract": "Multi-line imaging of the nearby disk galaxy NGC~1482 with the Taurus Tunable Filter (TTF) on the Anglo-Australian Telescope reveals a remarkable hourglass-shaped [N~II] $\\lambda$6583/H$\\alpha$ excitation structure suggestive of a galactic wind extending at least 1.5 kpc above and below the disk of the host galaxy. Long-slit spectroscopy confirms the existence of a large-scale outflow in this galaxy. The entrained wind material has [N~II] $\\lambda$6583/H$\\alpha$ ratios in excess of unity while the disk material is characterized by H~II region-like line ratios indicative of a starburst.  Expansion velocities of order 250 km s$^{-1}$ are detected in the wind material, and a lower limit of 2 $\\times$ 10$^{53}$ ergs is derived for the kinetic energy of the outflow based on the gas kinematics and the amount of ionized material entrained in the outflow. This is the first time to our knowledge that a galactic wind is discovered using excitation maps. This line ratio technique represents a promising new way to identify wind galaxy candidates before undergoing more time-consuming spectroscopic follow-ups. This method of selection may be particularly useful for samples of galaxies at moderate redshifts. ",
        "introduction": "In an effort to better constrain the morphology, kinematics, and origin of the warm ionized gas on the outskirts of galaxies we have obtained deep emission-line images of several nearby starburst and active galaxies using the Taurus Tunable Filter (TTF) on the 3.9m Anglo-Australian and 4.2m William Herschel Telescopes (see Veilleux 2001 for more detail).  In the course of this study, we discovered a remarkable emission-line structure in the early-type spiral galaxy NGC~1482. This galaxy has so far received relatively little attention in the literature. It is classified as a peculiar SA0/a in the Revised 3rd Catalogue (De Vaucouleurs et al. 1991) based primarily on the presence of a dust lane across the disk of the galaxy. Located at a distance of 19.6 Mpc (Tully 1988), it is an infrared-bright galaxy with log[L$_{\\rm IR}$/L$_\\odot$] = 10.5 (e.g., Soifer et al. 1989; Sanders, Scoville, \\& Soifer 1991), which is rich in molecular gas and dust (e.g., Sanders et al. 1991; Young et al. 1995; Chini et al. 1996) and is undergoing vigorous star formation (e.g., Moshir et al. 1990; Devereux \\& Hameed 1997; Thornley et al. 2000).  In a recent emission-line imaging survey of early-type spirals, Hameed \\& Devereux (1999) noticed the presence in NGC~1482 of ``filaments and/or chimneys of ionized gas extending perpendicular to the disk.''  The present study expands on the results of Hameed \\& Devereux, using deeper emission-line maps at H$\\alpha$ and [N~II] $\\lambda$6583 and complementary long-slit spectra. We find that the [N~II]/H$\\alpha$ excitation map of this galaxy is particularly useful at distinguishing between the star-forming disk and the entrained, shock-excited wind material.  This excitation signature could be used in the future to more efficiently identify powerful galactic winds in the local and distant universe. ",
        "conclusions": ""
    },
    "0201/astro-ph0201522_arXiv.txt": {
        "abstract": "In this paper I analyze the process of formation of thin current structures in the magnetosphere of a conducting accretion disk in response to the field-line twisting brought about by the rotation of the disk relative to the central star. I consider an axisymmetric force-free magnetically-linked star--disk configuration and investigate the expansion of the poloidal field lines and partial field-line opening caused by the differential rotation between the star and a nonuniformly-rotating disk. I present a simple analytical model that describes the asymptotic behavior of the field in the strong-expansion limit. I demonstrate the existence of a finite (of order one radian) critical twist angle, beyond which the poloidal field starts inflating very rapidly. If the relative star--disk twist is enhanced locally, in some finite part of the disk (which may be the case for a Keplerian disk that extends inward significantly closer to the central star than the corotation radius), then, as the twist is increased by a finite amount, the field approaches a partially-open configuration, with some field lines going out to infinity. Simultaneous with this partial field opening, a very thin, radially extended current layer forms, thus laying out a way towards reconnection in the disk magnetosphere. Reconnection, in turn, leads to a very interesting scenario for a quasi-periodic behavior of magnetically-linked star--disk systems with successive cycles of field inflation, opening, and reconnection. ",
        "introduction": "\\label{sec-intro}  Magnetic processes taking place in the magnetosphere above a thin accretion disk play an important role in establishing the structure of disk outflows (winds and jets), in regulating the accretion flow, and in angular momentum transfer. Of particular interest is the situation where there is a direct magnetic connection between the star and the disk (the so-called magnetically-linked star-disk system, see Fig.~\\ref{fig-geometry}). This connection may lead to a direct angular momentum exchange between the disk and the star and is thus relevant for neutron star spin-up/spin-down events (Ghosh \\& Lamb 1978, 1979; Wang 1987; Lovelace~et~al. 1995). It is also important because it provides a mechanism for direct channeling of accretion flow onto the polar regions of the star (Bertout, Basri, \\& Bouvier 1988; Lamb 1989; Patterson 1994; K{\\\"o}nigl 1991).  \\begin{figure}[tbp] \\begin{center} \\epsfig{file=geometry.eps,width=3.5 in} \\caption{Axisymmetric magnetically-linked star--disk system. \\label{fig-geometry}} \\end{center} \\end{figure}  The evolution and even the sole existence of such a configuration depends critically on the conducting properties of the disk, the star, and the overlying magnetosphere. This complex physical problem can be studied on various time scales; one should start, however, with the shortest relevant time scale, namely, the rotation period. On this time scale both the central star and the disk can usually be approximated by ideal conductors (with the exception of the case when the central object is a black hole), and so can the low-density magnetosphere (or corona) that lies above the disk. In this case there is no steady state because the differential rotation between the disk and the star leads to continuous twisting of the field lines. The magnetic field in the magnetosphere responds to this differential rotation  by rapid expansion driven by the increased toroidal field pressure. The field lines become elongated along a direction making a roughly $60^\\circ$ angle with the rotation axis. This expansion process has been studied in some detail in the framework of the force-free model.% \\footnote{The validity of the force-free approximation in these systems is justified by the fact that, due to very low density in the magnetosphere, the Alfv{\\'e}n velocity there greatly exceeds both the sound speed and the rotation speed.} In particular, semi-analytic self-similar models (van~Ballegooijen~1994, hereafter VB94; Lynden-Bell \\& Boily 1994, hereafter LBB94; Uzdensky et al. 2002a) have shown that the field-line expansion leads to an effective opening of the field lines after a finite (a few radians) twist angle (see Fig.~\\ref{fig-contour}). This finite-time singularity has also been observed in numerically constructed non-self-similar sequences of force-free equilibria (Uzdensky et al. 2002a) and also in non-force-free full-MHD numerical simulations by Goodson et al. (1997, 1999). This process is essentially very similar to the process of finite-time field-line opening of a twisted coronal field studied extensively in solar physics (see, e.g., Barnes \\& Sturrock 1972; Low 1986; Roumeliotis et al. 1994; Miki{\\'c} \\& Linker 1994; Wolfson 1995; Aly 1995). Thus, at present there is growing evidence that a finite-time opening of field lines in response to differential rotation is a generic feature of  force-free magnetospheres of magnetically-linked star--disk systems.  \\begin{figure}[tbp] \\begin{center} \\epsfig{file=contour.eps,width=5 in} \\caption{Field-line expansion caused by the relative rotation of the disk with respect to the star in the case of the $n=0.5$ self-similar model, where~$\\Psi_{\\rm disk}\\sim r^{-n}$ (This figure is taken from Uzdensky et al. 2002a). \\label{fig-contour}} \\end{center} \\end{figure}  It is, however, still not clear what happens after this opening. Basically, one can envision two possible outcomes: the field lines may either reconnect (across the separatrix between the two domains with the opposite direction of the radial field) and close back; or they may stay open indefinitely. In the first scenario (Aly \\& Kuijpers 1990, VB94, Goodson et al. 1997, 1999) some resistivity is present in the magnetosphere, and as a result, a small amount of flux gets immediately reconnected, leading to the formation of an X-point. As reconnection progresses, a toroidal plasmoid forms in the magnetosphere. This plasmoid contains most of the toroidal flux; it becomes completely detached from both the disk or the star and just floats away, presumably feeding a jet (Goodson et al. 1999). At the same time, the magnetic link between the star and the disk is reestablished, with the reconnected magnetic field lines being now twisted less than they were just before the reconnection event. The poloidal magnetic field tension then quickly contracts them back towards the star, and the system returns to its initial state. Thus, in this scenario the situation is manifestly time-dependent; the time evolution consists of periodic cycles of successive field line inflation due to the differential shearing, ``effective'' field-line opening, reconnection, and the relaxation of the reconnected magnetic field to the initial state (of lesser magnetic stress), which completes the cycle.  In the second scenario, investigated by Lovelace~et~al. (1995), there is no reconnection and a true steady state is established, whereby the field is at least partially open and the magnetic link between the star and a significant part of the disk is permanently broken.  The fact that the above two scenarios differ so dramatically raises the level of urgency of identifying the conditions under which each of them can be realized. The main difference between these two pictures is that {\\it magnetic reconnection} is allowed to take place in one but not in the other. Thus, as it often happens in astrophysics and space physics, the physics of magnetic reconnection plays a key role. This issue is by no means trivial and it is questionable whether it could be resolved even by a direct numerical simulation. Indeed, numerical simulations typically suffer from unrealistically large numerical resistivities that make it very easy for oppositely directed field lines to reconnect (as seems to be the case with the simulations by Goodson et al. 1999).  It is usually the case that in order to get efficient reconnection, a thin current sheet is needed. In the two-dimensional axisymmetric situation discussed here, a natural place where such a current sheet can arise as the system approaches the open-field state is the current concentration region between the two domains with oppositely directed radial magnetic field. While the field lines are still closed, the current is concentrated along the apex line~$\\theta=\\theta_{\\rm ap}(\\Psi)$.% \\footnote {Here, $\\theta_{\\rm ap}(\\Psi)$ marks the angular position of the apex of an inflated field line (an apex is defined as the farthest from the central star point on a field line). Usually, $\\theta_{\\rm ap}$ is close to $60^\\circ$ and depends only very weakly on the field line~$\\Psi$.} As the field opens up, this apex line turns itself into the separatrix between two regions of oppositely directed, open (i.e., extending to infinity) field lines that comprise the open-field configuration. If one now considers such a configuration (where the field lines have already been opened), then one discovers that the formation of a current sheet is essentially unavoidable, at least as long as the force-free approximation is valid.% \\footnote {As was pointed out by the referee of this paper, if a heated atmosphere is present above the disk, then, once the magnetic field opens up, the drop of magnetic intensity with distance may be fast enough for the disk wind far out to dominate and keep the field from reconnecting. The centrifugal wind could also hold the field open provided there is a lower, sufficiently heated atmosphere to feed that wind. This may be the implied reasoning behind the apparent neglect of the possibility of reconnection by Lovelace's et al. (1995).} Indeed, an {\\it open} magnetic field is potential, it has no toroidal field. This is because all toroidal flux created by the initial twisting has escaped to infinity as the result of magnetic field expansion and opening. The field becomes predominantly radial everywhere, with $B_r$ reversing across the separatrix $\\theta=\\theta_{\\rm ap}$. In the absence of the toroidal field $B_\\phi$ at $\\theta=\\theta_{\\rm ap}$, the pressure balance across the separatrix cannot be maintained, and this leads to the collapse of the magnetic configuration to one with an infinitesimally thin current sheet; the two oppositely directed magnetic fields move toward each other, forming a thin current layer. Finally, the collapse is stopped when resistive (or other non-ideal) effects become important in the layer. Because the current density is tremendously increased, magnetic reconnection can occur.  It is important to realize, however, that, strictly speaking, as long as one has a force-free magnetic field with the field lines {\\it closed}, one can never have a true, infinitesimally thin current sheet. Indeed, while the system is going through a sequence of equilibria of {\\it closed} field lines, i.e., for $t < t_c$ (where $t_c$ marks the moment time of field-line opening), there is always finite toroidal field, $B_\\phi$, present. This is because it is the toroidal flux that drives the expansion along $\\theta=\\theta_{\\rm ap}$, and in a force-free equilibrium its outward pressure balances against the poloidal field's large curvature. Thus, in a closed force-free configuration the toroidal flux is kept in place by the poloidal field tension and prevents poloidal field lines from contracting back to a less-stressed state. At the same time, this toroidal field provides the pressure in the $\\theta$~direction, which prevents the poloidal field collapse into a current sheet configuration.  One could ask, however, whether a current sheet can form {\\it asymptotically}, that is, whether the sequence of force-free equilibria, which governs the system's evolution, can asymptotically lead to stronger and stronger thinning of the current concentration region as one approaches the critical moment, so that the characteristic angular width of this region, $\\Delta\\theta$, goes to zero as $t\\rightarrow t_c$. Note that this does not always happen. For example, in the self-similar force-free model for a uniformly rotating disk (VB94; Uzdensky et al. 2002a), where all magnetic quantities are power laws of~$r$ with fixed power-law indices, $\\Delta\\theta$ does not go to zero but approaches a finite value that depends solely on the flux distribution $\\Psi_{\\rm disk}(r)$ on the disk surface. In particular, if $\\Psi_{\\rm disk}(r) \\sim r^{-n}$, then $\\Delta\\theta$ approaches a finite value of order a fraction of one radian if $n=O(1)$ (and is proportional to~$n$ in the limit $n\\rightarrow 0$), as the critical moment $t_c$ is approached.  Asymptotic current sheet formation was in fact observed in the self-similar model by LBB94 (in cylindrical geometry) and also in an essentially very similar work by Wolfson (1995) (in spherical geometry in the solar corona context). This is explained by the fact that the power law index~$n$ was forced to change during the evolution. In particular, in the cylindrical case considered by LBB94, all magnetic flux was required to go through the boundary at some small radius (compared with the radius under consideration). Therefore, for each value of the power exponent~$n$, the resulting equilibrium corresponded to a member of the VB94 family of solutions in the limit of infinite expansion, $t\\rightarrow t_c(n)$ (this limit was analyzed in detail by Uzdensky et al. 2002a). Since in the VB94 analysis $t_c(n)$ grows as $n\\rightarrow 0$, the gradual increase of the twist angle with time meant that, when considering the LBB94 solutions, one were bound to obtain a sequence of solutions with ever diminishing values of~$n$ [determined from the condition $t_c(n)=t$]. That is, in the LBB94 sequence of solutions, $n$ had to decrease and approached zero at some point, which lead to the reported thinning and formation of a current sheet. Related to this change of~$n$ was the fact that the field-line footpoints were allowed (and actually had) to move poloidally. If, instead, one insisted on having the footpoints tied firmly to the disk surface, then the flux distribution on this surface (which, in the self-similar model, must be a power law extending to arbitrary large radii ) could not change, at least on the short, rotation-period time scale considered here. Then $n$ and, hence, $\\Delta\\theta$ would both stay finite. Under these circumstances the prospects for magnetic reconnection to occur in a timely manner would be very slim, at least in the two-dimensional framework (see Uzdensky et al. 2002b).  Thus, the transformation of the current concentration region into a true, infinitesimally-thin current sheet, even in the asymptotic sense, cannot, in general, be taken for granted.  The failure of the self-similar force-free model to provide a plausible current-sheet formation (and hence reconnection) scenario makes it increasingly important to try to understand the asymptotic current-sheet formation process in a more general situation where the self-similar model does not apply. In particular, it is interesting to ask whether (still within a force-free framework) the current-concentration region can become very thin, thus indicating a way towards forming a true current sheet. If such a thin current layer does form, what is its structure? In particular, what is the $\\Psi$-profile of the angular width $\\Delta\\theta$ and how does it depend on the twist angle profile $\\Delta\\Phi(\\Psi)$?  It is important to realize that in the self-similar model there is no special radial scale and hence all the field lines must open up simultaneously. In a general, non-self-similar situation, however, there is a possibility of a {\\it partial field-line opening},% \\footnote {The importance of partial field-line opening in solar coronal processes has been recognized and emphasized by Low (1990) and by Wolfson \\& Low (1992), who suggested that a partially open field configuration may be energetically accessible even when a completely open configuration is not.} with the domain of open field lines being adjacent to that of closed field lines. As we shall discuss below, this can greatly facilitate the current-sheet formation process.  On the intuitive level, the basic idea of how a thin magnetic structure can form in the disk magnetosphere is simple. Imagine a {\\it non-uniformly} rotating disk. In particular, consider a system on the brink of the partial field-line opening: let there be a field line $\\Psi_c$ such that, as the system approaches a certain critical time $t_c$, the field lines outside of $\\Psi_c$ (i.e., $\\Psi<\\Psi_c$) tend to open, while those inside stay closed.\\footnote {We count poloidal magnetic flux $\\Psi$ from the outside inward, i.e., we set $\\Psi(\\pi/2,r\\rightarrow \\infty) \\rightarrow 0$.} Now, if $\\Delta\\Omega(\\Psi)$ [and hence the twist $\\Delta\\Phi(\\Psi)= \\Delta\\Omega t$] is nonuniform, say, rising outward, then the degree of expansion of field lines will also be non-uniform. Indeed, two neighboring field lines $\\Psi_1=\\Psi_c+\\delta\\Psi$ and $\\Psi_2=\\Psi_c-\\delta\\Psi$ will have their footpoints close to each other, while their apexes $r_{\\rm ap}(\\Psi_1)$ and $r_{\\rm ap}(\\Psi_2)$ will be very far apart since $r_{\\rm ap}(\\Psi_1)$ should stay finite whereas $r_{\\rm ap}(\\Psi_2) \\rightarrow \\infty$ as $t\\rightarrow t_c$. In other words, $d\\log r_{\\rm ap}(\\Psi)/d\\log\\Psi$ at $\\Psi=\\Psi_c$ will go to infinity as $t\\rightarrow t_c$. As we shall show in this paper, a thinning of the apex region (i.e., current layer formation) is characteristic for a situation like this, which is consistent with the spirit of the studies by LBB94 and by Wolfson (1995) who both show that a current sheet forms asymptotically when $n\\equiv d\\log\\Psi/d\\log r \\rightarrow 0$.  Thus, the main thrust of this paper is to demonstrate how twisting of field lines leads to the formation of thin current structures that asymptotically become thinner and thinner as the field approaches a partially-open state in finite time. In \\S~\\ref{sec-model} we present the basic geometry of the problem and a description of our model. In \\S\\S~\\ref{sec-twist}--\\ref{sec-radial-balance} we discuss the three main components of the model: \\S~\\ref{sec-twist} focuses on the purely geometrical relationship between the twist angle of the field lines and the relative strength of the toroidal field, while \\S\\S~\\ref{sec-theta-balance} and~\\ref{sec-radial-balance} describe our treatment of the~$\\theta-$ and the radial force balance equations, respectively. In \\S~\\ref{sec-analysis} we formulate and discuss the final set of differential equations describing the behavior of the angular thickness of the current concentration region. Also in that section we consider three specific examples: \\S~\\ref{subsec-const-u} deals with the constant-twist case, \\S~\\ref{subsec-critical-twist} investigates the behavior in the vicinity of the critical twist angle, and \\S~\\ref{subsec-numerical} describes a numerical solution of our equations for the case of a locally-enhanced twist angle. We present our conclusions in \\S~\\ref{sec-conclusions}.  Finally, we note that although we consider this problem in the accretion disk context, our methods and the main findings can be directly applied in the solar corona context. In particular, we note that the differential-rotation-driven process of partial field-line opening (and the detachment of the toroidal plasmoid associated with it) is essentially very similar to a coronal mass ejection; in addition, the partially opened field configuration considered in this paper is similar to a coronal helmet streamer configuration (see, for example, Low 2001). ",
        "conclusions": "\\label{sec-conclusions}  In this paper we addressed the issue of the formation of thin current structures in the magnetosphere of an accretion disk in response to the field-line twisting. We presented a simple analytical model that illustrates the asymptotic behavior of a force-free axisymmetric magnetic field above a thin conducting disk in the limit when such structures are about to form. We showed that there exist a finite critical twist angle, $\\Delta\\Phi_c$, such that the expansion of the poloidal field lines accelerates dramatically as $\\Delta\\Phi_c$ is exceeded. At the same time, in the case of a locally-enhanced twist,% \\footnote{ An important example of the locally-enhanced-twist case is a Keplerian disk whose inner edge is significantly closer to the central star than the corotation radius. In particular, if $r_{\\rm inner}<0.63\\, r_{\\rm co}$, then the relative star--disk angular velocity, $\\Delta\\Omega(r)=\\Omega_{\\rm disk}(r)-\\Omega_*$, reaches higher values in the localized inner region ($r_{\\rm inner}< r<r_{\\rm co}$) than at large distances ($r>r_{\\rm co}$).} a radially-extended thin conical current layer forms. Simultaneously, the field configuration approaches a partially-open state in finite time, with the apexes of a portion of the field lines going out to infinity.  The presence of a thin current sheet is usually considered a pre-requisite for reconnection of magnetic field lines. It is also believed that reconnection will occur whenever a thin current layer is formed. We therefore believe that magnetic field lines that have become effectively open as a result of their twist-driven inflation may subsequently reconnect via the thin current layer that has been formed simultaneously with their opening-up. This conclusion leads to important consequences regarding the time evolution of magnetospheres of magnetically-linked star--disk systems. In particular, we believe that, once the magnetic link between the star and the disk is broken by the partial field-line opening, a current layer will form along the separatrix between the stellar field lines and the disk field lines and, as a result, the link will be re-established through reconnection of these field lines. Subsequently, if reconnection process is sufficiently fast, the field lines will contract to a less-stressed configuration allowing a new cycle of twisting, inflation, opening, and reconnection to begin, as was suggested by Aly \\& Kuijpers (1990) and by VB94 and illustrated in numerical simulations by Goodson et al. (1999). This quasi-periodic scenario is characterized by very rich physics with non-steady and violent behavior, perhaps not too different from that of the solar corona (e.g., Low 2001). It can provide avenues for understanding such phenomena as disk winds, time-variable, knotted jets, episodes of rapid accretion, and variable accretion torque on the central star.  One should note, however, that a true opening will be preceded by the plasma inertia becoming important. The inertial effects will tend to retard the expansion, since the pressure of the toroidal field pressure will have to work not only against the poloidal field tension, but also against the plasma inertia. Hence the inertial effects will in effect act against the toroidal field removal and will not help to form a current sheet. At some point during the expansion they will need to be taken into account and, even in the absence of reconnection, a true finite-time field opening will not happen. Instead, one will have a transition from the force-free regime into the inertial regime (wind regime), which would require solving the full set of MHD equations, although probably under some simplifying assumptions.  I am very grateful to Stanislav Boldyrev, Arieh K{\\\"o}nigl, and Bob Rosner for very useful discussions and comments. I am also grateful to the anonymous referee for his or her deep and insightful criticism that was very beneficial for the paper. I would also like to acknowledge the support by the ASCI/Alliances Center for Astrophysical Thermonuclear Flashes at the University of Chicago under DOE subcontract B341495 and by the NSF grant NSF-PHY99-07949."
    },
    "0201/astro-ph0201536_arXiv.txt": {
        "abstract": "{We present pc-scale images of ten Compact Symmetric Objects (CSO) candidates observed with the European VLBI network (EVN). Five radio sources have been observed at 1.6 GHz, and five more at 2.3/8.4 GHz, the latter subsample with the inclusion in the VLBI array of 3 antennae normally used for geodesy. These objects were selected from existing samples of GHz Peaked Spectrum (GPS) radio sources with the purpose to find and/or confirm the CSO classification. These new VLBI observations allow us to confirm the classification of two CSO candidates, and to find a few new ones. The association of GPS radio galaxies with a CSO morphology is strengthened by our findings, and this result suggests an efficient way to increase  the number of known CSOs by means of VLBI observations of compact radio galaxies showing a convex radio spectrum. \\\\ ",
        "introduction": "Compact Symmetric Objects (CSOs) form a class of radio sources with distinctive radio properties. They are very powerful and compact sources with overall size $<$ 1 kpc, dominated by lobe/jet emission on both sides of the central engine, and are thought to be relatively free from beaming effects (Wilkinson et al. 1994). Their small size is likely due to their youth ($<$ 10$^4$ yr). This hypothesis has been largerly accepted after the detection and the estimate of the separation speed of the micro hot-spots in a couple of CSOs (Owsianik and Conway 1998, and Owsianik et al. 1998). Since then, expansion velocities of the outer edges have been detected or suspected in a handful of other CSOs (Fanti 2000 and references therein).  A unification scenario assumes that CSOs evolve into Medium-size Symmetric Objects (MSO), which, in turn, evolve into Large Symmetric Objects (LSO), i.e. large FRII radio sources (Fanti et al. 1995, Readhead et al. 1996ab, Snellen et al. 2000). The monitoring of the structural changes over time seems one of the more promising ways to understand the evolution of these objects, and a larger number of CSOs suitable for these repeated observations is desirable.  Most confirmed CSOs and CSO candidates are commonly found to have a global convex radio spectrum peaking at GHz frequencies, thus belonging to the class of the GHz Peaked Spectrum (GPS) radio sources.  These latter objects are known to show structures resolved only with VLBI. Only about 10\\% of them shows faint extended radio emission on scales of a few tens of kpc, while $\\sim$90\\% of them is entirely contained within the extent of the narrow-line region ($<$ 1 kpc) of the host galaxy (Stanghellini et al. 1998). GPS radio sources are usually identified with galaxies at low or intermediate redshifts (z$<$ 1), and quasars at higher redshifts (O'Dea et al. 1991, Snellen et al. 1999), and they make up a significant fraction  ($\\sim$10\\%) of the bright centimeter wavelength selected radio sources.  Many GPS radio sources have no optical counterpart yet, and these empty fields are most likely distant galaxies, too faint to be detected. Stanghellini et al. (1997a, 1998) find a strong correlation between the milliarcsecond (mas) morphology and the optical host in bright GPS sources. Galaxies are generally associated with CSOs while quasars have more often core-jet or complex morphology.  There are a few tens of sources in known GPS samples, and many more are in candidate lists; several of these sources have either no or poor VLBI images. In order to increase the number of known CSOs, we selected from the lists of known bright GPS samples (O'Dea et al. 1991, de Vries et al. 1997, Stanghellini et al. 1998, Dallacasa et al. 2000) the sources accessible to the EVN without enough structural information in the literature to permit a proper classification yet. Scheduling the radio sources to be observed in the allotted time we gave a higher priority to objects with known double morphology at a single frequency.   We report here the results of the two first observing runs. In a first session at 1.6 GHz with the EVN and MERLIN we observed 5 sources. At this frequency we could detect the presence of extended steep spectrum emission, determine the total angular size of the source, and possibly determine the low frequency spectral shape of the components, when other images at higher frequencies were available. Five further  radio sources were observed at 2.3/8.4 GHz with the EVN plus 3 antennae commonly used in geodesy experiments. These simultaneous observations at 2 frequencies are best suited to detect the cores. ",
        "conclusions": "We can summarize the results from our EVN observations at 1.66, 2.3 and 8.4 GHz for 10 GPS radio sources as follows:  i) we confirm the CSO nature of 1345+125 revealing a component which is very likely to be the core.  ii) for the radio source 1404+286 we confirm from ground based VLBI observations the findings by Kameno et al. (2000) i.e. that the radio spectrum is affected by FFA rather than SSA, or maybe both. This is an important observational clue because it implies the presence of a dense ionized ambient medium somewhere between us and the strongly absorbed radio emitting region.  iii) we also found more CSO candidates which need need further observations to be confirmed.  The relationship between radio mas morphology and optical host is rather respected in the small sample considered here, although also quasars may appear in the CSO candidate list (e.g. 1518+047). Excluding 1751+278 which does not show a convex radio spectrum and it has been erroneously classified as a GPS radio source, we find that 5 galaxies (we include J1335+5844 in the galaxy group as empty fields are mostly identified with weak galaxies once sensitive optical observations become available) are confirmed or candidate CSO, one quasar is also a CSO candidate and 2 quasars are unresolved or slightly resolved objects in our images. Therefore the observations of  GPS radio galaxies seem to be an efficient way to find CSOs.  The MERLIN data at 1.66 GHz do not reveal any extended emission in four of the five sources observed. Only 1345+125 is very sligthly resolved, in agreement with the size seen in the EVN image.  Since the cores are known to have highly inverted spectra, multi-frequency and especially high frequency sensitive observations are needed to confirm the CSO nature for the new CSO candidates discovered here. If they are confirmed they will be very suitable objects to monitor for separation speed, due to their small size (thus possibly faster separation velocity) and the presence of strong and compact components (hot-spots?) at their edges."
    },
    "0201/astro-ph0201470_arXiv.txt": {
        "abstract": "{Observed self-gravitating systems reveal often fragmented non-equilibrium structures that feature characteristic long-range correlations. However, models accounting for non-linear structure growth are not always consistent with observations and a better understanding of self-gravitating $N$-body systems appears necessary. Because unstable gravitating systems are sensitive to non-gravitational perturbations we study the effect of different dissipative factors as well as different small and large scale boundary conditions on idealized $N$-body systems. We find, in the interval of negative specific heat, equilibrium properties differing from theoretical predictions made for gravo-thermal systems, substantiating the importance of microscopic physics and the lack of consistent theoretical tools to describe self-gravitating gas. Also, in the interval of negative specific heat, yet outside of equilibrium, unforced systems fragment and establish transient long-range correlations. The strength of these correlations depends on the degree of granularity, suggesting to make the resolution of mass and force coherent.  Finally, persistent correlations appear in model systems subject to an energy flow. ",
        "introduction": "Most astrophysical structures result of gravitational instabilities, from large scale cosmological structures down to planets.  Yet, among the least understood topics in astrophysics we find galaxy formation and star formation, which both involve fragmentation and the nonlinear growth of structures occurring during the non-linear phases of gravitational instability.  Perhaps, one of the fundamental reason why fragmentation and structure formation via gravitational instability appears so difficult is that we lack of consistent theoretical tools allowing to combine gravity with gas physics. Indeed, too often ignored is that classical thermodynamics does not hold for gravitating systems, because these are non-extensive in the thermodynamical sense (Landsberg \\cite{Landsberg72},\\cite{Landsberg84}; Tsallis \\cite{Tsallis99}; Plastino \\& Plastino \\cite{Plastino99}).  Actually, many other natural systems do not respect the requisites of thermodynamics. Such systems often feature interesting phenomena such as growing long range correlations or phase transitions.  Among the symptoms of a fundamental deep problem in gravitating systems is the appearance of negative specific heat (Lynden-Bell \\& Lynden-Bell \\cite{Lynden77}; Lynden-Bell \\cite{Lynden98}), which was seen for a long time as a paradox in statistical mechanics, since negative specific heat was thought to be impossible.  Presently, the only available approach to follow the nonlinear phases of gravitational instabilities is to carry out numerical simulations. Among all the existing methods, $N$-body techniques are thought to be the most effective to simulate self-gravitating systems as well the continuous case as the granular phases.  Yet, despite the considerable success of these methods in reproducing many observed features, many fundamental problems remain.  As mentioned above, the fragmentation and structure formation is not clearly understood.  Related to this, CDM simulations conflict with observations at galactic scales (Moore \\cite{Moore99}, James et al. \\cite{Bullock01}, Bolatto et al. \\cite{Bolatto02}), and no theory of the ISM is presently able to {\\it predict\\/} the conditions of star formation.  Most of the time the star formation process relies on recipes with little physical constraint.  In situations were $N$-body simulations have success (e.g. hot stellar systems) gravitational dynamics is sufficient to account for their main global properties, additional microscopic physics can be neglected. But when gravitational instability via fragmentation involves small scale physics, the outcome may be strongly dependent on the properties of the small scale physics.  In other words, in situations where the growth on singularities triggered by gravity is allowed, the chaotic nature of gravitating systems make them sensitive to the perturbations induced by non-gravitational physics.  Therefore it is important to understand the properties of $N$-body systems subjected to various perturbations. For these purposes, a numerical study of perturbed, self-gravitating $N$-body systems is carried out.  Among the relevant perturbations we expect that boundary conditions at small and large scales, as well as dissipative factors can play a key role.  In order to characterize the individual effects of perturbations, in the tradition of analytical models, one is advised to deliberately use simplified models.  A study of dissipative systems is important because such systems may develop long-range correlations. In the typical ISM, radiative cooling is very effective and induces a temporary energy-flow leading the system far from equilibrium (Dyson \\& Williams \\cite{Dyson97}).  From laboratory experiments it is well known that systems outside of equilibrium may spontaneously develop spatio-temporal structures (Glansdorff \\& Prigogine \\cite{Glansdorff71}; Nicolis \\& Prigogine \\cite{Nicolis77}; Prigogine \\cite{Prigogine80}; Melo \\cite{Melo94}).  A permanent energy-flow is induced when energy loss due to dissipation is replenished, that is, when the system is continuously driven, e.g., by time-dependent boundary conditions. Such systems may develop persistent long-range correlations.  Astrophysical examples for this are, the growth of structures in cosmological simulations, or the long term persistence of filamentary structures in shearing flows (Toomre \\& Kalnajs \\cite{Toomre91}; Huber \\& Pfenniger \\cite{Huber01a}; Wisdom \\& Tremaine \\cite{Wisdom88}; Salo \\cite{Salo95}; Pfenniger \\cite{Pfenniger98}). Among other things, the effect of time-dependent potential perturbations on dissipative self-gravitating spheres is studied in this paper.  In the next Section we briefly review some theoretical results of the thermodynamics of self-gravitating isothermal spheres. The model is presented in Sect. 3 and the applied methods to carry out the structure analysis for the simulated systems are explained in Sect. 4. The results are presented in Sect. 5-7.  In Sect. 5 quasi-equilibrium states of $N$-body models are compared with analytical predictions. Sect. 6 is dedicated to a study of long-range correlations appearing in the interval of negative specific heat during the collapsing transition of gravitating systems. Finally, in Sect. 7 the evolution of systems subjected to an energy-flow is discussed. ",
        "conclusions": "First, equilibrium states of $N$-body models were compared with analytical models. Subsequently, the findings resulting from this comparison are summarized:  \\begin{itemize}  \\item On the one hand, equilibrium properties of $N$-body models agree with predictions made by analytical models. An example is the energy interval of negative specific heat. One the other hand, discrepancies were found, such as the way the collapsing phase transition, separating a high-energy homogeneous phase from a low-energy collapsed phase, develops in the interval of negative specific heat. These discrepancies suggest: 1.) Small scale physics becomes relevant for the system evolution when the growth of singularities triggered by gravitational instabilities is allowed. 2.) Analytical models based on the Gibbs-Boltzmann entropy are not strictly applicable to non-extensive self-gravitating systems. Yet, not all of the equilibrium properties found by maximizing the Gibbs-Boltzmann entropy are expected to change if a fully consistent, generalized thermostatistical theory  is applied.  \\end{itemize}  \\noindent Second, the collapsing transition was studied in systems with strong dissipation. The findings are:  \\begin{itemize}  \\item Dissipative self-gravitating systems develop outside of equilibrium, in the interval of negative specific heat, transient long-range correlations. That is, fragmentation and nonequilibrium velocity-dispersion-size relations, with striking resemblance to those observed in the ISM, appear during the collapsing transition, when the dissipation time is shorter than the dynamical time. This suggests that nonequilibrium structures in self-gravitating interstellar gas are {\\it dynamical and highly transient}.  \\item Besides the dissipation strength and the initial noise, the granularity turns out to be a crucial parameter for the strength of the resulting long-range correlations, substantiating the importance of a coherent mass and force resolution. That is, phase-space correlations are stronger in the fluid limit than in a granular phase. The opposite holds for spatial correlations.  This substantiates The inverse behavior of fragmentation strength and phase-space correlation strength is found in all simulations and is typical for self-gravitating systems.  \\end{itemize}  \\noindent Finally, systems subject to a permanent energy-flow were studied. We find:  \\begin{itemize}  \\item Typically driven dissipative systems evolve to a high-energy homogeneous phase or undergo a mono-collapse. Yet, model systems with a local energy dissipation can develop persistent phase-space correlations, but a persistent, hierarchical fragmented structure is not observed. This suggests, that matter that has passed through a collapsing transition has to be replenished at large scales in order to maintain a hierarchical structure at molecular cloud scales.  \\end{itemize}"
    },
    "0201/astro-ph0201193_arXiv.txt": {
        "abstract": "We use simulations of the formation and evolution of the galaxy population in the Local Universe to address the issue of whether the standard theoretical model succeeds in producing empty regions as large and as dark as the observed nearby ones. We follow the formation of galaxies in a \\lcdm universe and work with mock catalogues which can resolve the morphology of LMC sized galaxies, and the luminosity of objects 6 times fainter. We look for a void signature in sets of virialized haloes selected by mass, as well as in mock galaxy samples selected according to observationally relevant quantities, like luminosity, colour, or morphology. We find several void regions with diameter $10 \\hMpc$ in the simulation where gravity seems to have swept away even the smallest haloes we were able to track. We probe the environment density of the various populations and compute luminosity functions for galaxies residing in underdense, mean density and overdense regions. We also use nearest neighbour statistics to check possible void populations, taking $L_{*}$ spirals as reference neighbours. Down to our resolution limits, we find that all types of galaxies avoid the same regions, and that no class appears to populate the voids defined by the bright galaxies. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0201/astro-ph0201120_arXiv.txt": {
        "abstract": "{ Retrograde waves with frequencies much lower than the rotation frequency become trapped in the solar radiative interior. The eigenfunctions of the compressible, nonadiabatic, Rossby-like modes ($\\epsilon$-mechanism and radiative losses taken into account) are obtained by an asymptotic method assuming a very small latitudinal gradient of rotation, without an arbitrary choice of other free parameters. An integral dispersion relation for the complex  eigenfrequencies is derived as a solution of the boundary value problem. The discovered resonant cavity modes (called $R$-modes) are fundamentally different from the known $r$-modes: their frequencies are functions of the solar interior structure, and the reason for their existence is not related to geometrical effects. The most unstable $R$-modes are those with periods of $\\approx$\\,1--3\\,yr, 18--30\\,yr, and 1500--20\\,000\\,yr; these three separate period ranges are known from solar and geophysical data. The growing times of those modes which are unstable with respect to the $\\epsilon$-mechanism are $\\approx 10^2, 10^3,$ and $10^5$ years, respectively. The amplitudes of the $R$-modes are growing towards the center of the Sun. We discuss some prospects to develop the theory of $R$-modes as a driver of the dynamics in the convective zone which could explain, e.g., observed short-term fluctuations of rotation, a control of the solar magnetic cycle, and abrupt changes of terrestrial climate in the past. % ",
        "introduction": "The 22-year magnetic cycle of solar activity is the most prominent phenomenon of several large-scale dynamic events that occur in the Sun. (Really, the magnetic half cycles or sunspot number cycles vary in length between 9--13 years, and 11 yr is an average of the $\\approx20$ half-cycles available.) An explanation of the basic mechanism underlying this fameous phenomenon is the fundamental challenge of solar physics. The achievements of the theory of the $\\alpha$-$\\omega$ dynamo turned out to be a great success. However, neither all observations of magnetic and flow fields nor the radiation fluxes which are related to this phenomenon and which are measured at the surface of the Sun or indirectly, by helioseismology, in its interior, can be explained unambigously in this way. Although our present work is not directly related to the dynamo theory, we will outline here those difficulties which have common points with our results.   \\subsection{Some problems of dynamo theory }.   As a consequence of our imperfect knowledge of basic characteristics of turbulent convection as well as meridional circulation and details of the rotation of the Sun's interior, the solutions of the dynamo equations become functions of many free, unknown parameters (e.g. Stix 1976). For instance, by clever combinations of these parameters it is possible to get from kinematic theory an oscillatory magnetic field with a 22-year period and a growing amplitude. However, another choice of these parameters leads to waves of growing amplitude for other periods. So one could draw a butterfly diagram not only with an 11-year periodicity. It remains still an open question which of the clever combinations resulting in a solar-like 22-year activity cycle is realized in the Sun. We could not find a work on the dynamo wave problem, showing that just the 22-year period is preferred among others with a maximum growth rate and with the spatial scales required for solar activity. Instead, many authors pointed out that the cycle period of 22 years is hard to explain (Stix 1991; Gilman 1992; Levy 1992; Schmitt 1993; Brandenburg 1994; Weiss 1994; R\\\"udiger \\& Arlt 2000).  From the solution of the inverse problem of helioseismology (e.g. Tomczyk et al. 1995) it is known that the convective envelope of the Sun is rotating with a latitude dependence of the angular velocity similar to that of the surface but almost rigid in radial direction. A stronger radial gradient which is required for the $\\alpha$-$\\omega$ dynamo mechanism is located in a shallow layer (thickness $\\approx 0.05 R_{\\odot}$ (Kosovichev 1996), where $R_{\\odot}$ is the solar radius) immediately below the convective zone --- the tachocline (Spiegel \\& Zahn 1992). Below the tachocline up to a depth of at least $0.5 R_{\\odot}$ the radiative interior is rotating with an angular velocity law similar to that of a solid-body. The  question arises: what compels the Sun to rotate in such a strange manner, which is different from the generally accepted, theoretically predicted stable rotation law? How to handle a dynamo theory for which the `$\\omega$' area is separated from the `$\\alpha$' area over a large part of the extent of the convective zone? In order to solve this problem Parker (1993) has put forward the idea of an interface dynamo, the basic features of which existed already in earlier dynamo models (Steenbeck et al. 1966). To close the cycle of such a stretched dynamo it is necessary to have some mechanism delivering toroidal magnetic flux, arising by the shear of differential axisymmetric rotation (Cowling 1953) in the tachocline, to the `$\\alpha$' dynamo area (e.g. Moffatt 1978; Krause \\& R\\\"adler 1980). To get a solar-like magnetic activity it is necessary to suppose the existence of a huge ($\\approx 10^5$\\,G) toroidal magnetic field to create enough magnetic buouancy for the leakage of magnetic flux and to solve the tilt problem of lifting loops (e.g. Caligari et al. 1998). Moreover, a high magnetic diffusivity contrast between the convective envelope and the underlying radiative core should be assumed to solve the quenching problem of the $\\alpha$ effect (see, e.g., Fan et al. 1993; Cattaneo \\& Hughes 1996). However, it is a major challenge for any dynamo model to produce such strong fields.  The idea of the interface dynamo was further developed, e.g. by Charbonneau \\& MacGregor (1997). Later, a fit to the real solar rotation profile with its latitudinal and radial dependencies has been included by Markiel \\& Thomas (1999), but so far no satisfactory solar-like oscillatory solutions for the interface dynamo have been found. Growing wave solutions are suppressed by the latitudinal shear.   \\subsection{Spinning-down rotation problem}.   Mechanisms for braking the solar internal rotation are also under discussion. The  character of the core rotation is not clear because here the accuracy of helioseismic inversions gets worse (Chaplin et al. 1999) and the results seem to be in contradiction with the oblateness measurements (Paterno et al. 1996). There are some suggestions that a deceleration of the radiative interior depends on the transport of angular momentum between this region and the convective zone. For instance, Mestel \\& Weiss (1987) supposed that even a weak large-scale magnetic field would be sufficient to couple very efficiently the interior and the  convective zone, leading essentially to solid body rotation. In this way the magnetic torques can also extract angular momentum from the  radiative interior (e.g. Charbonneau \\& MacGregor 1993).  The wave mechanism for the solution of this problem is more popular. Schatzman (1993), Zahn et al. (1997), and Kumar \\& Quataert (1997) have concluded that the solid rotation of the radiative interior is a direct consequence of the effect of internal gravity waves. Gravity waves generated near the interface between the convective and radiative regions transport retrograde angular momentum into the interior, thereby spinning it down. Here the main idea is that the isotropically generated gravity waves become anisotropic due to Doppler shifts of frequencies in the differentially rotating Sun. In that way for anisotropic retrograde and prograde waves the radiative damping is different, and the residual negative angular wave momentum may compel the solar radiative interior to co-rotate with the convective zone. This idea has been further developed by Kumar et al. (1999) including a toroidal magnetic field to explain the existence of the unstable shear layer `tachocline'. However, Ringot (1998) has shown that a quasi-solid rotation of the radiative interior cannot be a direct consequence of the action of internal gravity waves produced in the convective zone. Gough (1997) questioned this idea emphasizing that the mechanism can work only if the waves are generated with strong amplitudes to transport the required angular momentum. This means, resonance waves are required, but such waves may penetrate only to distances less than $10^{-5}R_{\\odot}$ beneath the convective zone due to the strong radiative damping. These waves must deposit their angular momentum before returning to the convective zone, but not before penetrating far into the radiative interior.  For the wave mechanism the question of an anisotropic propagation relative to the azimuthal rotation is a key moment. Fritts et al. (1998) have shown that convection, penetrating into the stratified and strongly sheared tachocline, can produce preferentially propagating gravity waves.  There have also been speculations that the rotation of the core may be variable, perhaps with a time scale of the solar cycle (e.g. Gough 1985). The present paper is along these lines.  From our short discussion we conclude that the convective envelope and the radiative interior are coupled to each other through a certain global agent, resulting in almost co-rotation. To advance the solution of the problem the dynamo theory should take into account the presence of this global agent. We suppose that really this agent is provided by waves with the following properties:  Waves should represent large-scale global eigenoscillations of the Sun. Their origin must be related to rotation, they must be strongly anisotropic with respect to the azimuthal angle. Looking at the characteristics of the solar cycle we immediately see the high coherency of these global motions (the constant periods, phase shifts, amplitudes, the latitude appearence, etc.). Activity grows in the first phase with a timescale which is considerably shorter than the decay time in the second phase; this fact and the quick eruptive release of energy by the reconnection mechanism indicate that the waves must be unstable.  It is noteworthy that the inner gravity waves do not fulfill these requirements. The quesion is whether the $r$-modes do?   \\subsection{$r$-modes}.   In a non-rotating sphere ($\\Omega=0$, where $\\Omega$ is the angular frequency of solar rotation) the wave motion is subdivided into two non-coupling components: spheroidal $p-$, $f-$ and $g-$modes (for which the main restoring forces are pressure gradient and buouancy) and toroidal modes (e.g. Unno et al. 1989). Toroidal modes are degenerated horizontal eddy motions confined to a spherical surface with a radius $r$ for which $\\omega=0$, $\\dvrg{\\vec{v}}=0$, and $\\vec{v}=Q^m_l(r)\\times (0,\\frac{1}{\\sin\\theta}\\pd{}{\\phi},-\\pd{}{\\theta}) Y^m_l(\\theta,\\phi)$. Here $Y^m_l$ is the spherical harmonic with a degree $l$ and order $m$, $\\theta$ is the colatitude, $\\phi$ is the azimuthal angle in the spherical polar coordinates, $\\vec{v}$ is the fluid velocity field, $\\omega$ is the angular frequency of the fluid motion, and $Q^m_l(r) $ is an arbitrary amplitude function. Toroidal modes have zero radial velocity but have non-zero radial vorticity, $(\\rot{\\vec{v}})_r\\ne 0$ (for the spheroidal modes it is vice versa). These modes do not alter the equilibrium configuration.  When a slow rotation ($\\Omega^2<\\Omega^2_g=GM/R^3$) is included the spheroidal modes are slightly modified but they keep their main properties. Degeneracy of toroidal modes is removed only partially by the rotation, and quasi-toroidal waves -- known as {\\it r-modes} -- appear with a non-zero frequency of $\\omega\\approx 2\\Omega m/l(l+1)$ in the rotating frame (Papaloizou \\& Pringle 1978; Brayn 1889). Usually the governing equations of the $r$-modes are obtained by expanding the initial physical variables of the equations in the rotating system into power series with respect to the small parameter $(\\Omega/\\Omega_g)^2$ ($\\approx 10^{-4.7}$ for the Sun, e.g. Papaloizou \\& Pringle 1978; Provost et al. 1981; Smeyers et al. 1981; Saio 1982). These power series practically describe the deviation of the surface of the star from its initial spherical state, resulting from rotation through Coriolis and centrifugal forces. As a result of the deformation of the spherical surface with a radius $r$ the radial vorticity of the toroidal modes cause a surface pressure perturbation through the Coriolis force. However, the $r$-modes practically keep the main properties of toroidal flows: $v_r\\approx 0$, $\\dvrg{ \\vec{v}} \\approx 0$. The degeneracy of the $r$-modes is that their frequencies hardly depend on $Q^m_l(r)$, i.e. they are independent from the inner structure of the star. For the $l=1$ modes the frequency in the inertial system is again close to zero, $\\omega\\approx 0$ (Papaloizou \\& Pringle 1978). The $r$-mode equations define the amplitudes $Q^m_l(r)$, and taking into account the next terms with small $\\Omega^2/\\Omega^2_g$ in the series practically does not change the frequencies.  Due to the fact that the $r$-modes are practically surface deformation waves, some similarity of these waves to the surface gravity waves or to the $f$-modes is apparent.  For high $l$ the $f$-modes are an analogy to surface gravity waves in a plane-parallel fluid with $\\omega^2=gk_{\\perp}$. In the Cowling approximation $f$-modes with $l=1$ have zero frequency too, $\\omega\\approx 0$ (Unno et al. 1989). This corresponds to a parallel displacement of the whole star. For high $l$ the $f$-mode frequencies are also independent of the inner structure, with $\\omega^2\\approx l\\Omega^2_g$ (Gough 1980). So, $r$-modes are also fundamental rotating modes with an inertial frequency, $\\omega\\le 2\\Omega$.  For the Sun the properties of $r$-modes have been investigated in great detail by Wolff et al. (1986) and Wolff (1998; 2000; and refs. therein).  Some properties of the $r$-modes are also similar to those of the Rossby waves in geophysics (Pedlosky 1982). Similar to the Rossby waves and unlike the $g$-modes the $r$-modes are strongly anisotropic. They propagate only in azimuthal direction, opposite to rotation (i.e. they are retrograde waves in the co-rotating frame). Because we are interested in length scales corresponding to those of large sunspots, we have to consider $r$-modes with $l\\approx 100$. To get oscillations with periods of years ($\\omega/\\Omega\\approx 10^{-2}$) we must choose $m\\approx l\\gg 1$, just such $r$-modes are physically more interesting (Lockitch \\& Friedman 1999). However, in the case of high $l$ the amplitudes of the $r$-modes will be concentrated near the surface of the Sun (Provost et al. 1981; Wolff 2000), and so they can actively interact with convective motions (Wolff 1997; 2000). Because for these modes $v_r\\approx 0$ and $\\dvrg{\\ \\vec{v}}\\approx 0$, their chance to take part in the redistribution of angular rotation momentum in the radiative interior is low. Note that the slow solar differential rotation does not change the behavior of such $r$-modes with $m=l\\gg 1$ (Wolff 1998).  Looking for further analogies between waves connected with gravity and with rotation, we remember that beside the surface gravity waves there exist internal gravity waves with $\\omega^2\\approx N^2k^2_x/(k_x^2+k_y^2)$, the frequencies of which depend on the inner structure ($N$ is the Brunt-V\\\"ais\\\"al\\\"a frequency). Similar to these waves there exist `true' Rossby (not deformation) waves, the frequency of which depends also on the internal structure.    \\subsection{Rossby waves}   We include here a short review on the main features of Rossby waves; they have been investigated in great detail in geophysics (e.g. Pedlosky 1982; Gill 1982). In the simplest case, that is in a plane-parallel, homogeneous, rotating layer, the dispersion relation for the Rossby waves is $\\omega\\approx 2\\Omega\\beta k_x/(k_x^2+k_y^2+ k_z^2)$. Here $k_x$ is the wave number perpendicular to the rotation axis, $k_z$ is expressed by the internal deformation radius of Rossby which depends on the Brunt-V\\\"ais\\\"al\\\"a frequency, $\\beta$ is the transverse gradient (in $y$ direction) of the Coriolis parameter: a vertical component of the `planetary' vorticity $2\\vec\\Omega$ in the given local point. Unlike the $r$-modes the Rossby wave frequencies are functions of the internal  structure and have maximum dependence on the gradient $\\beta$: $\\omega\\to 0$ if $k_x\\to 0$ and if $k_x\\to\\infty$.  Any disturbance of the local flow in a rotating frame may generate waves of the Rossby type. These waves exist only if there is a gradient of the potential vorticity $\\vec\\Pi =(\\vec{\\omega}_a\\cdot\\nabla\\Xi)/\\rho$. Here an absolute vorticity is the sum of the relative and the planetary vorticities, $\\vec{\\omega}_a=\\rot{\\vec{v}}+2\\vec\\Omega$, $\\Xi$ is any conserved scalar quantity, $\\der{\\Xi}{t}=0$ (for instance, for adiabatic motion that could be the entropy or the density in the case of incompressible plasma). The Rossby wave motion is a solution of the nonlinear equation for transport of $\\Pi$. The potential vorticity is conserved if the medium is barotropic ($\\nabla\\rho\\times\\nabla p=0$) and if there are no torques. The rotation of the frame is added to any vorticity in the velocity field. Any motion within a rotating fluid serves as a potential source for vorticity.  The relative vorticity may be evoked by the geometrical surface as well as by internal gradients. It depends on the choice of the function $\\Xi(r,\\theta,\\phi)$ and on $\\vec{\\Omega}(r,\\theta)$. For example, an unevenness of the ocean bottom causes the topographic Rossby waves, or a dependence of the Coriolis parameter on the earth latitude ($F=2\\Omega\\sin\\varphi$, where $\\varphi$ is the geographic latitude) is the main cause of atmospheric Rossby waves.  In the solar dynamo context the ability of Rossby waves to induce solar-like magnetic fields has been considered by Gilman (1969). Here the mechanism for sustaining the Rossby waves is a latitudinal temperature gradient in a thin, rotating, incompressible convective zone.  To interpret the dynamical features of large-scale magnetic fields the Rossby vorticies excited within a thin layer beneath the convective zone are considered by Tikhomolov \\& Mordvinov (1996) as the result of a deformation of the lower boundary of the convective zone.   \\subsection{$R$-modes}   From the discussion in Subsection 1.4 we conclude that just Rossby-like waves could be suitable for our requirements. As the main driving mechanism  we choose a latitudinal (or horizontal) differential rotation, $\\Omega=\\Omega(\\theta)$. Baker \\& Kippenhahn (1959) have pointed out that the uniform rotation of a star is not a typical case. Low frequencies (periods of years) could easily be obtained searching for the eigenoscillations of the Sun's radiative interior, where the gradient of the rotation speed is close to zero (in accordance with the helioseismology results). Large scales such as those associated with sunspots ($k_xR_{\\odot}\\sim 100$) decrease the frequencies too. Similar to the $r$-modes the Rossby waves are strongly anisotropic (retrograde waves), but unlike the $r$-modes these waves are concentrated close to the solar center. These results have already been obtained by Oraevsky \\& Dzhalilov (1997), who investigated the trapping of adiabatic, incompressible Rossby-like waves in the solar interior. In the present work we take into account compressibility for the nonadiabatic waves. We look for unstable waves. It is clear that the necessary condition for the Kelvin-Helmholtz shear  instability, $4N^{2}-(r\\der{\\Omega}{r})^{2}<0$ (Ando 1985), is not fulfilled in the radiative interior. Then we decided to include the thermal $\\en$-mechanism of instability which is favoured at low frequencies (Unno et al. 1989). To balance the $\\en$-mechanism the radiative losses in the diffusion regime are included. To exclude all geometrical effects we ignore the influence of the spherical surface at the given radius. The dispersion relation in the limit of adiabatic incompressiblity and at very low frequencies is the same as that for Rossby waves in geophysics. In order to distinguish these rotational body waves from the $r$-modes we call them {\\it R-modes} (Rossby rotation).  The governing fourth order equation is obtained from the basic equations in Sect.\\,2. Some qualitative analysis of the wave cavity trapping is done for the simpler adiabatic case in Sect.\\,3. Using the asymptotic solutions obtained in Sect.\\,4 the complex boundary value problem is solved in Sect.\\,5. The calculation of the eigenfrequencies and the instability analysis are done in Sect.\\,6. The obtained unstable modes are shortly discussed in Sect.\\,7. ",
        "conclusions": "In the present paper we have shown that toroidal eddy flows which are degenerated in a non-rotating fluid can become a reservoir of various branches of oscillatory modes when the degeneracy is removed by rotation. The mechanism depends on the condition for the existence and alteration of the relative vorticity as well as on the stellar rotation rate and its gradient. Apparently at least for slowly rotating stars ($\\Omega<\\Omega_g$) the rotation waves could be divided into two types: $r$-modes with high frequencies ($\\omega\\le\\Omega$) which are independent from the inner structure and mainly caused by geometrical effects, and the $R$-modes with low frequencies ($\\omega \\ll \\Omega$) which depend on the inner structure and are considered in the present paper. This classification is similar to that of $f$- and $g$- modes or to that of surface and body tube modes of magnetic cylinders. Note that the properties of Coriolis forces and  ponderomotive forces in MHD are very similar to each other. Both rotation modes are prototypes of the geophysical Rossby waves.  We investigated the instability problem of the $R$-modes sustained by a very small latitudinal gradient of the rotation rate in the solar radiative interior. Among the eigenoscillations three modes with periods of $\\approx$\\,1--3\\,yr, 18--30\\,yr, and 1500--20\\,000\\,yr turn out to be maximum unstable to the $\\en$-mechanism. Here the smoothing effect is the radiative damping. All of these instabilities are in the range of high radial node numbers $n$ which indicates that the applicability of the asymptotic solution is satisfied.  The 22-yr modes with a growing time of $\\approx1000$~yr are of particular interest with respect to the solar activity cycle problem. In the simpler case when adiabatic $R$-modes are considered in an incompressible fluid, $\\sigma_n$ in Eq.~(\\ref{dis3}) is independent of the wave number and of the frequency for very low frequencies. Then in the azimuthal direction the phase and  group velocities are \\be v_{px}\\approx \\frac{\\beta}{\\bar{k}^2_{\\perp}+\\sigma^2_n}, \\ \\ v_{gx}\\approx -\\beta\\frac{\\bar{k}^2_x-\\bar{k}^2_y-\\sigma^2_n} {(\\bar{k}^2_{\\perp}+\\sigma^2_n)^2} , \\ee respectively, where the velocities are normalized to $\\Omega_{\\odot}R_{\\odot}$ $ \\approx 1.6$\\,km/s. In our case $\\beta>0$, i.e. the angular velocity is decreasing towards the pole, similar to the behavior at the solar surface. The $x$-axis in our coordinate system is directed opposite to the direction of rotation, $v_{px}>0$ and $v_{gx}<0$, moreover, $k_y^2\\approx -k_x^2$ and for finite $n$ we have $\\bar{k}_x^2>\\sigma_n^2$ (see Fig.2). Then the retrograde $R$-modes transport energy along the direction of rotation. Our treatment in a Cartesian coordinate system does not allow to determine the direction of energy transport by wave packets relative to the equator in the meridian plane. The estimate of $k_y^2\\approx -k_x^2$ is crude, and to determine the exact dependence on $k_y=k_y(k_x,\\omega,\\beta)$ the 2D boundary value problem must be solved.  A nice property of the Rossby waves is that every monochromatic mode is a solution of the full nonlinear hydrodynamic equations. It means, that we should expect the development of nonlinear $R$-modes with large amplitudes. We could also expect that just in this nonlinear regime the toroidal magnetic flux will be lifted from the upper boundary of the cavity (the tachocline) to the surface. The energy release of the nonlinear waves could be accomplished by magnetic reconnection. Here it is possible that toroidal currents are generated via a twist of toroidal magnetic field lines by the cyclonic flows of regular $R$-modes with fixed characteristics. Parker (1955) as well as Steenbeck et al. (1966; see also Krause \\& R\\\"adler 1980) have suggested for the dynamo process that such a mechanism, the $\\alpha-$effect, is working by turbulent motions under the influence of Coriolis forces.  Our present model points out the possibility of forced oscillations instead of a self-excited dynamo to solve the solar cycle problem, and this with the correct period of 22 yr. Similar ideas are due to Tikhomolov (2001) who has recently suggested a hydrodynamic driving of the 11-yr sunspot cycle. We expect that in our model --- contrary to classical dynamo models --- a huge toroidal magnetic field of $\\approx 10^5\\,G$ will no longer be required to explain the buoyant rise of magnetic flux tubes appearing at the surface with small tilt angles and at low latitudes: the external nonzero upflow produced by the regular vortical $R$-modes could trigger the eruption of stable magnetic flux tubes stored in the overshoot region. There is still a smaller peak of the growth rate (Fig.5) at 100 years; such a period is observed as a modulation of the 11/22 yr cycles.  There is observational evidence for the short-period oscillations as well: From helioseismic sounding Howe et al. (2000, 2001) have recently discovered variations of solar rotation with a period of 1.3 yr in the lower convective zone. Quasi-two year modes are very likely seen regularly in various solar data (e.g. Waldmeier 1973; Akioka et al. 1987; Rivin \\& Obridko 1992). The existence of two magnetic cycles (the main 22-yr and the quasi-biennial period) on the Sun has been reported by Benevolenskaya (1996; 1998). So far the origin of these modes was not yet clear. Terrestrial quasi-biennial oscillations have been clearly seen in tropical meteorological radiosonde data, and a possible solar origin by related phenomena in the solar interior, Rossby waves in particular, has been discussed as well (McIntire 1994).  The long-period oscillations in the broad range $1.5\\times 10^3$\\,--\\,$2.0\\times 10^4$~yr, with a maximum growth rate around 4500~yr, could be the cause of abrupt changes of the global terrestrial climate in the past: Dansgaard-Oeschger events, these are abrupt onsets of warm periods during the last ice age, had mean distances of 4500~yr, but they were distributed over a larger period range, similar to that in our model, with shortest distances often around 1500~years (see, e.g., Ganopolski \\& Rahmstorf 2001). These events were caused by changes of the thermohaline circulation of the ocean, which in its turn were probably triggered by changes in the solar energy output."
    },
    "0201/astro-ph0201316_arXiv.txt": {
        "abstract": "The hypothesis that radiation of Sgr A* is caused by accretion onto a supermassive compact object without the events horizon is studied. The main equations of the accretion and a relativistic equation of the trasfer radiation are obtained. The synchrotron spectrum in the vicinity of the maximum is considered. ",
        "introduction": "An analysis of stars motion in the dynamic center of the Galaxy gives evidences for the existence of a supermassive ($2.6\\cdot10^{6}M_{\\odot}$ ) compact object .( See reviews \\cite{Goldwurm}, \\cite{MeliaFalke01}). There are three kinds of explanation of the observed perculiarity of the object:  1.The gas accretion onto the central object - a supermassive black hole.  2. The ejection of the magnetized plasma from the vicinity of the Schwarzshild radius of the above object.  3. The explanations are based on some hypotheses about another nature of the central objects (neutrino ball, boson stars).  In the present paper we begin an investigation of the assumption that the radiation of Sgr A* is conditioned by a spherically symmetric accretion onto a supermassive compact object without the events horizon. The possibility of the existence of such a kind of objects follows \\cite{Verozub95}, \\cite{VerKoch00} from our gravitation equations \\cite{Verozub91}$.$ This is a stable configuration of the Fermi-gas with radius $R$ less than the Schwarzshild radius $r_{g}$ ( $R$ is about $0.04\\ r_{g}$ for the mass $2.6\\cdot 10^{6}M_{\\odot}$) which unlike a black hole has no the events horizon.  It was shown earlier \\cite{VerBan98} that the accretion onto objects of this kind does not contradict the low observed bolometric luminosity of Sgr A* ($<$ $10^{37}$ $erg/s$ ).  The spectrum of the radiation near its maximum ( $10^{13}\\div10^{14}$ $Hz$ ), that supposedly comes from the vicinity of the object surface, is obtained. It follows from the results that the analyzed model can be considered as one of the explanation of Sgr A* radiation. ",
        "conclusions": "To clear up question about the nature of Sgr A* is issue of the day since the problem is the nature of the supermassive compact objects. Since 1992 a number of models have been created that are able to explain Sgr A* spectrum. It is the possibility of an ambiguous explanation of the spectrum that means that we need for more rigorous methods of the calculation of the Sgr A* spectrum. It is possible only by a simultaneous solution of the relativistic hydrodynamics and transfer equations.  I am very grateful to S.Zane for a kind explanation of their approach to the transfer equation and J. Schmid-Burg for sending me his paper on this subject.."
    },
    "0201/astro-ph0201066_arXiv.txt": {
        "abstract": "We explore a mechanism for the formation of the first globular clusters, operating during the assembly of dwarf galaxies at high redshifts, $z\\ga 10$. The substructure in the dark matter and the corresponding potential wells are responsible for setting the cluster scale of $\\sim 10^{5}M_{\\odot}$. The second mass scale in the formation problem, the stellar scale of $\\sim 1 M_{\\odot}$, is determined in turn by the processes that cool the gas. We address the origin of the first, cluster mass scale by means of three-dimensional numerical simulations of the collapsing dark matter and gaseous components. We find that the gas falls into the deepest dark subhalos, resulting in a system of $\\sim 5$ proto-globular clouds. The incipient globular clusters lose their individual dark halos in the process of violent relaxation, leading to the build-up of the general dark halo around the dwarf galaxy. ",
        "introduction": "The origin of globular clusters (GCs) is a longstanding challenge in astrophysics, ever since Peebles \\& Dicke (1968) have tried to link their formation to the conditions in the early universe, briefly after the epoch of recombination (at redshift $z\\sim 1000$). Despite considerable progress in our theoretical understanding (e.g., Fall \\& Rees 1985; Ashman 1990; Kang et al. 1990; Murray \\& Lin 1992; McLaughlin \\& Pudritz 1996; Padoan, Jimenez, \\& Jones 1997; Nakasato, Mori, \\& Nomoto 2000; Ashman \\& Zepf 2001; Cen 2001), the fundamental question of why it is that at early cosmological times bound aggregates of $\\sim 10^{5}$ stars were able to form, remains unsolved. Two recent developments, however, have significantly improved the prospects for real progress.  Observations with the {\\it Hubble Space Telescope} have revealed several merging or recently merged systems which contain extremely luminous young stellar aggregates (e.g., Whitmore \\& Schweizer 1995). These systems have estimated sizes and masses close to those of Galactic GCs. Therefore, in extreme environments like starburst galaxies, GCs might still be able to form in the present-day universe where we can directly probe the formation process. The second recent development is the increase in computational power which enables us to address the complex physics of collapsing and fragmenting gas with high resolution and the addition of the important physical processes. Therefore, the renewed investigation of globular cluster formation is very timely, and holds the promise of understanding star formation on its grandest scale.  It is likely that there are many avenues leading to the formation of GCs (e.g., Ashman \\& Zepf 1998). In this {\\it Letter}, we explore one of them and ask: {\\it Could the first globular clusters have formed during the initial stages in the hierarchical build-up of cosmic structure?} Within popular variants of the `cold dark matter' (CDM) model, the first dwarf galaxies of mass $\\sim 10^{8}M_{\\odot}$ are expected to collapse at $z \\ga 10$. Forming GCs in these dwarf systems might provide an answer to the old puzzle of how to simultaneously account for the two characteristic mass scales involved in the problem: the cluster scale of $\\sim 10^{5} M_{\\odot}$, and the stellar scale of $\\sim 1M_{\\odot}$, respectively. In the context of our model, the cluster scale is set by the substructure in the dark matter (DM) component, providing the potential wells in which the proto-globular clouds are assembled. The second, stellar, mass scale is then determined by the cooling physics of the star forming gas. We here address the origin of the cluster mass scale by performing numerical simulations of the collapsing DM and gas components. The emergence of the stellar mass scale will be investigated in subsequent, higher-resolution work.  The possible connection between GC formation and DM subhalos has previously been explored semi-analytically by Peebles (1984), Rosenblatt, Faber, \\& Blumenthal (1988), and C\\^{o}t\\'{e}, West, \\& Marzke (2001). Following the demonstration that GCs do not have individual dark halos (Moore 1996), this scenario had lost much of its initial appeal. Considering GC formation during the early stages of the CDM bottom-up hierarchy, however, might also provide a way out of this problem. Due to the flatness of the CDM spectrum on the smallest scales, there is a strong `cross-talk' behavior,i.e., the roughly simultaneous collapse of all scales (e.g., Blumenthal et al. 1984). The incipient GCs, provided they have condensed sufficiently, could therefore lose their individual dark halo in the process of violent relaxation without being disrupted themselves. ",
        "conclusions": "The collapse and virialization of the dwarf galaxy results in the formation of six high density clumps, with masses $10^{5}\\la M \\la 10^{7}M_{\\odot}$, and typical radii $\\sim 10$ pc. The resulting clump masses are determined by the mass spectrum of the DM substructure, provided the gas can cool sufficiently to fall into the shallow DM potential wells. This is the case in the presence of an efficient coolant that is able to operate at $T\\la 10^{4}$K, but is prohibited after the universe underwent reionization. In our simulations, the highly condensed gas clouds have lost their individual DM subhalos, which we tentatively ascribe to the exceptionally strong tidal forces acting during the relaxation of a dwarf galaxy at high redshift.   The expected metallicity distribution of clusters formed in this way needs further investigation. We note that the time between turnaround and virialization, which marks the interval over which GCs are assembled in our model, is around $10^{8}$yr, and thus in principle there is time for the nucleosynthetic products from one cluster to reach the gas destined to form another cluster. It is currently unclear whether this could explain the significant spread in metallicity observed among the clusters in the Fornax and Sagittarius dwarfs (e.g., Buonanno et al. 1998), or whether these data require that some of the clusters form at significantly later epochs."
    },
    "0201/astro-ph0201250_arXiv.txt": {
        "abstract": "Observations indicate that much of the interstellar gas in merging galaxies may settle into extended gaseous disks.  Here, I present simulations of disk formation in mergers of gas-rich galaxies.  Up to half of the total gas settles into embedded disks; the most massive instances result from encounters in which both galaxies are inclined to the orbital plane.  These disks are often warped, many have rather complex kinematics, and roughly a quarter have counter-rotating or otherwise decoupled central components.  Disks typically grow from the inside out; infall from tidal tails may continue disk formation over long periods of time. ",
        "introduction": "Gas falling into the nuclei of interacting and merging galaxies has attracted much attention.  Nuclear inflows fuel the bursts of star formation responsible for the anomalous colors (Toomre \\& Toomre 1972; Larson \\& Tinsley 1978) and extraordinary IR luminosities (Joseph \\& Wright 1985; Sanders \\& Mirabel 1996) of merging galaxies; the stars formed in this manner may play a crucial role in transforming merger remnants into elliptical galaxies (Kormendy \\& Sanders 1992). Numerical simulations show that these inflows occur when gravitational torques remove angular momentum from shocked galactic gas (Noguchi 1988; Hernquist 1989; Combes, Dupraz, \\& Gerin 1990; Barnes \\& Hernquist 1991, 1996).  In typical cases $\\sim 60$ percent of the gas initially distributed throughout the progenitor disks may wind up in a nuclear cloud with dimensions of $\\sim 0.1 {\\rm\\,kpc}$.  Less attention has been paid to the gas which {\\it fails\\/} to reach the nuclei of merger remnants.  Extended H{\\small I}, H{\\small II}, and X-ray emission is often seen in merging galaxies, but the gas responsible is usually too dilute to directly fuel violent star formation or AGN activity.  There has been little numerical work on the fate of extended gas in mergers, in part because it's expensive to continue a simulation once a compact nuclear gas cloud has formed.  There are reasons to think that extra-nuclear gas may play a role in the evolution of merger remnants.  The well-studied merger remnant NGC~7252 exhibits extended emission from ionized gas with remarkably complex kinematics (Schweizer 1982, 1998): an inner disk with radius $\\sim 3 {\\rm\\,kpc}$ and a well-defined axis of rotation is surrounded by a region which shows rapid minor-axis rotation and an apparent velocity reversal.  This inner disk, seen as a vivid spiral in HST images (Whitmore et al. 1993), contains $\\sim 10^{9.5} {\\rm\\,M_\\odot}$ of molecular and ionized gas (Dupraz et al. 1990).  A comparable amount of atomic gas lingers in the tidal tails; these return gas to the main body of the remnant at at rate of $\\sim 4 {\\rm\\,M_\\odot\\,yr^{-1}}$ (Hibbard et al. 1994; Hibbard \\& Mihos 1995). Whatever the fate of this gas, it's worth asking how a system with such complex kinematics could form in the first place.  In this paper I use simulations to explore the possibility that extended gas disks and associated structures can form as part of the merger process.  Each simulation begins with a pair of disk galaxies falling together on a parabolic relative orbit, and follows their collision and the eventual formation and initial relaxation of a merger remnant.  The interstellar material is modeled by including a component obeying the equations of motion of an isothermal gas.  This approach is deliberately simple-minded when compared to simulations which include star formation, feedback, and other processes (e.g.~Katz 1992; Mihos \\& Hernquist 1994; Gerritsen \\& Icke 1997; Springel 2000). Star formation is clearly important for the evolution and eventual fate of merging galaxies, but there's no evidence that stellar processes are essential for the formation of {\\it gas\\/} disks like the one in NGC~7252.  My intention is to present a simple model which appears capable of explaining the disks observed in merger remnants. ",
        "conclusions": "Gas disks like those described here were briefly reported in earlier simulations of dissipative mergers between disk galaxies (Hernquist \\& Barnes 1991; Barnes \\& Hernquist 1996).  These studies used the same general methodology as the present work, but differed in several specifics.  First of all, in earlier work the thermal evolution of the gas was explicitly computed, including radiative cooling with a cut-off at $T = 10^4 {\\rm\\,K}$.  This had little consequence for the gas in regions of moderate to high density where cooling times are short; however, the gas tails in the present experiments are broader and smoother since they are not cooled by adiabatic expansion. Second, the earlier studies used King (1966) distributions for the bulges and halos of the initial galaxy models.  As a result, these models may have been more susceptible to bar instabilities (Mihos \\& Hernquist 1996), and their relative orbits decayed appreciably faster (Barnes 1998).  Third, the earlier simulations used the TREESPH code (Hernquist \\& Katz 1989); the code used in the new simulations borrows several features from TREESPH, but it adopts a different smoothing prescription for gravitational forces, a more efficient and robust tree algorithm, a fixed number of neighbors in SPH summations, and a different time-step algorithm.  Despite using somewhat different models and codes, these studies all yield broadly similar results; the process of disk formation appears to be insensitive to such details.  Simulations of polar ring formation in galaxy mergers have produced remnants with very large rings or annuli (Bekki 1998).  These rings arise from a relatively restricted set of initial conditions which don't overlap with the ones considered here.  Disk formation is also seen in simulations of hierarchical clustering which include gas as well as cold dark matter (e.g.~Evrard, Summers, \\& Davis 1994; Navarro \\& Steinmetz 1997; Dom\\'{\\i}nguez-Tenreiro, Tissera, \\& S\\'aiz 1998).  These disks are generally identified with {\\it spiral\\/} galaxies, and they form largely from pristine gas which was not previously part of another galaxy.  In common with the disks seen in the present experiments, disks formed by hierarchical clustering generally grow from the inside out.  \\subsection{Origins of counter-rotation}  A rather striking result from this study is the large number of remnants with counter-rotating or otherwise misaligned nuclei.  In some cases, counter-rotation is easily explained; for example, it's no great surprise that the RETrograde encounters all produce remnants with counter-rotating components.  Likewise, some misaligned components appear to contain gas originating almost exclusively from one of the two progenitors, and it's likely that such ``unblended'' structures retain some memory of the orientation of their parent disks.  But the origin of the misalignment is not always so evident.  The nuclear disk in remnant POL~1:1~C is a case in point.  In terms of overall structure this remnant is very similar to one previously described (Hernquist \\& Barnes 1991); in both a fairly massive nuclear gas disk counter-rotates with respect to an outer gas disk and stellar component.  This resemblance may not seem surprising, since the same initial disk geometry ($i_1 = 71^\\circ$, $\\omega_1 = 90^\\circ$, $i_2 = -109^\\circ$, $\\omega_2 = 90^\\circ$) was used in both experiments.  But the initial orbit used in the previous study was {\\it wider\\/}, with a first pericenter at $r_{\\rm peri} = 0.4$.  Of the encounters in the present study, the one with initial conditions most closely matching those used previously is POL~1:1~D, and POL~1:1~D does {\\it not\\/} yield a counter-rotating nucleus.  It's not clear why this encounter doesn't produce a counter-rotating nucleus like the one seen in the earlier study, while it's close relative POL~1:1~C provides a very good match.  One possibility is that the halos used in the present study are a bit less effective at promoting orbital decay; as a result, the merging nuclei arrive at their final encounter with slightly more orbital angular momentum than they had in the earlier experiment.  The mechanism which forms counter-rotating nucleus in objects like remnant POL~1:1~C remains obscure.  The nucleus of this remnant is already counter-rotating when it forms, and it's difficult to determine when and how this material first acquired its rotation; angular momentum can only be measured with respect to a center, and before the nucleus forms the system has no unique center.  This ambiguity in defining the center could actually be part of the puzzle -- material co-rotating with respect to the center as defined in one way may be counter-rotating with respect to the center defined in another.  Large-scale hydrodynamic interactions seem to be important; many of the counter-rotating or misaligned nuclei in the present sample form when galaxy pairs arrive at their second passage still dressed in extensive gas disks.  A better understanding of the roles of gravitational and hydrodynamic torques in the formation of counter-rotating nuclei may require calculations with higher spatial resolution than available at present.  \\subsection{Observational connections}  The results presented here invite comparisons with observations of gas disks in merger remnants.  As mentioned in the introduction, NGC~7252 is an obvious example; the central disk of ionized and molecular gas in this system is about the same size as the disks produced in these experiments, and the peculiar kinematics of this disk may be due to a strong warp of the kind seen in many of these simulations.  In addition, H{\\small I} velocities confirm that gas is falling back into the remnant from the tidal tails (Hibbard et al.~1994).  NGC~7252 has been modeled as the result of a close and fairly direct encounter between two disk galaxies of comparable mass (Hibbard \\& Mihos 1995). This model nicely reproduced the morphology and kinematics of the tails, but the calculations did not include a dissipative component which could form a central disk.  It's worth repeating this calculation with a combined N-body/SPH code; a single model reproducing both the tails {\\it and\\/} the central disk would be a fairly impressive accomplishment, and might help constrain the gas content of NGC~7252's progenitors.  The peculiar elliptical galaxy NGC~3656 (Balcells 1997, Balcells et al. 2001) may be a second example.  This galaxy contains shells and a pair of faint tails which suggest a merger of two disk galaxies.  It also has a star-forming dust lane which corresponds to an extended disk of H{\\small I}.  This disk is visibly warped, and the outer edge of the disk is kinematically contiguous with gas at larger radii which may be falling in from the tidal tails.  The estimated gas accretion rate in NGC~3656 is an order of magnitude lower than in NGC~7252; this may imply that NGC~3656 is at a later stage in its evolution than NGC~7252.  NGC~5128, an elliptical galaxy with an active nucleus and an extended disk of dust and gas, may represent an even later stage in the evolution of gas-rich merger remnants.  Although the warped disk in this system has been interpreted as the remains of an accreted gas-rich satellite (e.g.~Malin, Quinn, \\& Graham 1983), there are some grounds to suspect that this galaxy is the result of a fairly major merger (e.g.~Schweizer 1998).  These include the ripples visible in deep optical images (Malin et al.~1983), the H{\\small I} fragments surrounding the galaxy (Schiminovich et al.~1994), and the misaligned rotation revealed by kinematics of planetary nebulae (Hui et al.~1995).  {\\it Nuclear\\/} gas disks and rings seem to be very common in merging galaxies selected by infrared luminosity (e.g.~Downes \\& Solomon 1998).  In remnants with single nuclei the gas tends to rotate more rapidly than the stars, indicating some degree of kinematic decoupling (Genzel et al. 2001), though it's not clear if the extreme decoupling in remnant POL~1:1~C has yet been seen in the observations.  Disks in systems which appear to still have {\\it two\\/} nuclei are not so easily explained.  Molecular-line observations of Arp~220 and NGC~6240 have been interpreted in terms of a rotating disk located {\\it between\\/} the nuclei (Scoville, Yun, \\& Bryant 1997; Tacconi et al.~1999; Tecza et al.~2000).  In the simulations, gas driven inward before the galaxies merge always accumulates in disks or bars around the individual nuclei; moreover, an incipient disk between the nuclei would be torn apart by gravitational fields as the orbits of the nuclei decay.  An inter-nuclear disk might be stable if it was more massive than either nucleus (Tecza et al.~2000), but the amount of gas required seems quite extravagant, and earlier reports of a peak in the stellar velocity dispersion between the nuclei of NGC~6240 (Lester \\& Gaffney 1994; Doyon et al.~1994) are not supported by recent HST observations (Gerssen et al.~2001).  On the other hand, the interpretation of Arp~220's nuclei as bright spots in ``a warped molecular gas disk'' (Eckart \\& Downes 2001; Downes \\& Solomon 1998) seems entirely consistent with the present numerical results.  \\subsection{Galaxy transformation}  Simulations indicate that extended gas disks can form in mergers of spiral galaxies.  These disks contain between $20$ and~$60$ percent of the total gas in the initial galaxies, and may extend to several times the remnant half-light radii.  It's worth noting that the initial galaxy models used in these experiments were not particularly gas-rich, and that the gas started with the same distribution as the disk stars.  In real spiral galaxies the atomic gas is more extended than the stars; mergers of such systems should form remnants with even larger and more massive gas disks.  Moreover, the gas content of disk galaxies is generally expected to increase with redshift; remnants with disks containing $\\sim 10$ to~$15$ percent of their luminous mass would very likely result by doubling the gas fraction in the initial galaxy models.  These disks, if subsequently converted to stars, would be fairly hard to detect photometrically (Rix \\& White 1990) unless viewed from a favorable orientation (e.g.~Scorza \\& Bender 1990).  On the other hand, disks could significantly influence the observed {\\it kinematics\\/} of merger remnants.  Disks are kinematically cold, so stars formed in such disks would add narrow features to the stellar absorption-line spectra of early-type galaxies.  Moreover, disks rotate faster than pressure-supported components, so these narrow features will be systematically offset from the broader profiles due to the rest of the galaxy.  The combined velocity profile of a disk plus spheroid thus appears asymmetric, with a steep prograde wing and a shallow retrograde wing (e.g.~Franx \\& Illingworth 1988; Bender 1990; Rix \\& White 1992; van der Marel \\& Franx 1993).  Such profiles actually appear to be fairly typical in elliptical galaxies with measurable rotation (Bender, Saglia, \\& Gerhard 1994).  The observed velocity profiles may be compared with those predicted for remnants of purely dissipationless mergers between disk galaxies (Bendo \\& Barnes 2000, Naab \\& Burkert 2001).  As a rule, simulated equal-mass mergers produce remnants with the {\\it opposite\\/} velocity profile asymmetry -- that is, they have shallow prograde wings and steep retrograde wings.  For unequal-mass mergers the situation is not so clear -- one study finds profiles similar to those produced in equal-mass mergers, while the other reports cases which exhibit the same sense of asymmetry seen in observations of early-type galaxies. Naab \\& Burkert (2001) suggested that the formation of extended disks is {\\it necessary\\/} if mergers are to account for the observed kinematics of elliptical galaxies, but worried that the gas might not retain enough angular momentum to build extended disks.  Given the contradictory results on unequal-mass mergers, it may be premature to insist that all mergers must form disks.  On the other hand, the present work indicates that disks of the requisite size can form quite easily in both equal-mass and unequal-mass mergers.  If the above explanation for line profiles in elliptical galaxies is correct, one might expect different spectral lines to show different kinematic signatures.  The lines produced by the younger stars making up a disk should be narrower than, and systematically offset from, the lines produced by the older stars in the spheroid.  This might be tested with high signal-to-noise spectra obtained using large optical telescopes."
    },
    "0201/astro-ph0201299_arXiv.txt": {
        "abstract": "If gamma-ray bursts trace the cosmic star formation rate to large redshifts, their prompt and delayed emissions provide new tools for early universe cosmology. In addition to probing the intervening matter via absorption lines in the optical band, GRB continua also contribute to the evolving cosmic radiation background. We discuss the contribution of GRBs to the high-energy background, and the effect pair creation off low-energy background photons has on their observable TeV spectra. ",
        "introduction": "Cosmic gamma-ray bursts (GRBs) have redshifts comparable to or perhaps even larger than those of quasars. Indeed, they are the most energetic explosions in the universe, with energies (uncorrected for beaming) of order M$_\\odot$c$^2$. Their host galaxies are often sub-L$_*$, but actively forming stars at rates typical for galaxies in the early universe. The current paradigm associates GRBs with the formation of black holes in massive, rapidly rotating stars, or with the merger of compact star binaries. GRBs thus trace directly, or perhaps with a short delay, the cosmic star formation history, and may be the most easily detectable signposts of the first generation of stars (e.g., [1] and references therein). The redshifted gamma-ray flux from GRBs contributes to the evolving radiation background of the universe, as discussed in the next section, and at the same time serves as a probe of the cosmic radiation field through electron-positron pair creation absorption of their highest energy photons [2][3].  The cosmic microwave background (CMB) provides abundant soft photons for pair production of very high energy photons (in the TeV - PeV regime), but at lower energies (GeV-TeV regime) the target photons are optical and IR photons produced by stars and reprocessed by surrounding dust. Cosmic chemical evolution is intimately linked to the cosmic star formation history, and the present day extragalactic background light (EBL) provides a record of that history. Gamma-ray sources, such as GRBs and blazars, probe the evolution of this photon field through absorption effects at high energies (e.g., [2][3]). There are only three nearby active galaxies for which this absorption effect has been observed, Mrk 421, Mrk 501 (both at z = 0.03), and BL Lac (at z = 0.044). TeV emission from GRBs has only been reported for GRB970417a [4]. GRB power spectra ($\\nu$f$_\\nu$) typically peak at photon energies of a few hundred keV, but their power-law high energy emision may extend well into the GeV or even TeV regime. EGRET aboard the Compton Observatory has established that emission above 100 MeV is common, and in the case of GRB940217 a maximum photon energy of E $\\sim$ 20 GeV was determined [5].  Theoretical models (e.g., [6]) certainly suggest that GeV-TeV emission should be expected for a significant fraction of all bursts. The next generation GLAST experiment is expected to observe a large number of GRBs with spectral coverage up to 300 GeV. Ongoing improvements of ground-based experiments (VERITAS, HESS, HEGRA, MILAGRO, MAGIC, ...) lead to reduced sensitivities and thresholds, thus overlapping with space-based experiments. It will thus be possible to explore GRB spectra from the X-ray regime to the TeV regime, and the effects of propagation effects such as the above mentioned electron-positron pair creation must be taken into account.  To correct for $\\gamma\\gamma$ absorption, it is necessary to determine the cosmic evolution of the target photon distribution function, which we refer to as the metagalactic radiation field (MRF). In the third section of this paper we briefly describe our simulations of the evolving low-energy MRF, and demonstrate the extinction effect in the high energy part of GRB spectra. The gamma-ray horizon of the universe can perhaps be probed with GRBs, which would provide another powerful tool for the study of stellar evolution on the cosmic scale. GRB detections point to the onset of star formation in the universe, and their high energy spectra probe the production of light throughout the cosmic ages. ",
        "conclusions": ""
    },
    "0201/astro-ph0201145_arXiv.txt": {
        "abstract": "We report on the measurement of the two-point correlation function and the pairwise peculiar velocity of galaxies in the IRAS PSCz survey. The real space two-point correlation function can be fitted to a power law $\\xi(r) = (r_0/r)^{\\gamma}$ with $\\gamma=1.69$ and $r_0=3.70 \\mpc$. The pairwise peculiar velocity dispersion $\\sigma_{12}(r_p)$ is close to $400 \\kms$ at $r_p=3\\mpc$ and decreases to about $150 \\kms$ at $r_p \\approx 0.2 \\mpc$.  These values are significantly lower than those obtained from the Las Campanas Redshift Survey, but agree very well with the results of blue galaxies reported by the SDSS team later on. We have constructed mock samples from N-body simulations with a cluster-weighted bias and from the theoretically constructed GIF catalog. We find that the two-point correlation function of the mock galaxies can be brought into agreemnt with the observed result, but the model does not reduce the velocity dispersions of galaxies to the level measured in the PSCz data. Thus we conclude that the peculiar velocity dispersions of the PSCz galaxies require a biasing model which substantially reduces the peculiar velocity dispersion on small scales relative to their spatial clustering. The results imply that either the cosmogony model needs to be revised or the velocity bias is important for the velocity dispersion of the IRAS galaxies. ",
        "introduction": "\\begin{figure} \\plottwo{f1a.eps}{f1b.eps} \\caption{ The predictions of CDM models vs the observation for the projected two-point correlation function ({\\it left}) and the pairwise velocity dispersion ({\\it right}). Triangles show the observational results. The mean value and the $1\\sigma$ limits predicted by the cluster-weighted bias model are shown by the thick and thin lines respectively, and the mean values of the GIF simulation and the SCDM CLW model by the dot-dashed and dashed lines (without error bars). The SCDM curve of the correlation function is shifted vertically by a factor of $1/\\sigma_8^2$ to account for the necessary linear bias in this model.  } \\label{fig1}\\end{figure}  We measure the two-point correlation function and the pairwise velociity dispersion for the PSCz galaxies following exactly the procedure of JMB98. The projected 2PCF $\\wrp{}$ are presented in Figure~1 with the filled triangles. The data can be well fitted by a power law of $\\xi(r) = (r_0/r)^\\gamma$ with $r_0 = 3.70 \\mpc$ and $\\gamma =1.69$.  The slope is shallower and the amplitude is lower than those for the LCRS (JMB98), reflecting the fact that the PSCz galaxies are preferentially in the field environments.  The error bars are not plotted for clarity but are very close to the error bars predicted by the mock samples of the LCDM model with $\\Omega_0=0.3$, $\\lambda_0=0.7$ and $\\sigma_8=1.0$ (the dotted lines).  For the mock samples, we have applied the CLW bias with $\\alpha=-0.25$, which is found to be in good agreement with the PSCz data (the solid line); even the wiggly structure of the PSCz $\\wrp{}$ below $0.8 \\mpc$ is recovered.  To follow up this point a bit more, we have also constructed mock catalogs from the GIF simulation data (Kauffmann et al. 1999a,b).  From the simulated catalogue, we select those galaxies with $\\Delta V_{bg} \\equiv V_{b}-V_{g} \\ge 1$, where $V_{b}$ and $V_{g}$ denote the V-band magnitudes of the bulge and the whole galaxy.  This encompasses about 80 percent of the galaxies in the GIF catalog. The resulting $\\wrp{}$ again fits the observations quite well (Fig.~1) for $r_p \\simgt 1\\mpc$. While the amplitude of $\\wrp{}$ for the GIF simulation is a bit larger than the PSCz data for $r_p \\simlt 1\\mpc$, it still lies close to the $+1 \\sigma$ error line of our CLW mock samples.  So these semi-analytic models of galaxy formation which incorporate physical processes like star formation and supernova explosions in some global way have a similar effect of reducing the number of galaxies per unit dark matter mass as our simple bias prescription.  Figure~1 also displays the pairwise velocity dispersion for the self-similar infall model. The PVD is much lower than that for the LCRS: $\\sigma_{12}$ just reaches $300 \\kms$ at $r_p = 1 \\mpc$, whereas $\\sigma_{12}(1 \\mpc) = 570 \\pm 80 \\kms$ for the LCRS (JMB98). This is again qualitatively consistent with the fact that spirals have smaller random motions than the galaxies that reside in big clusters. Comparing with the PVD for the CLW and GIF mock samples, we find that the model predictions in all cases significantly exceed the estimate for the PSCz.  Only around $r_p \\approx 3 \\mpc$ the disagreement is not serious, especially if the large error bars are taken into account.  In fact, we may speculate that a CDM model with $\\Omega_0 = 0.2$ may even produce quite a good fit to the data at larger $r_p$, since the amplitude of the PVD scales with $\\sigma_8 \\Omega_0^{0.6}$. Reducing the values of $\\sigma_{12}$ accordingly, brings agreement on scales larger than $3 \\mpc$. There is, however, no way to reproduce the steep decrease towards small values ($\\sim 150 \\kms$) at $r_p = 0.2 \\mpc$ of the PSCz data. The pairwise velocity dispersions are rather similar in the SCDM ($\\Omega_0=1$ and $\\sigma_8=0.6$) and in the LCDM models. ",
        "conclusions": "We have analyzed the data set of the IRAS PSCz galaxies (Saunders et al. 2000), and computed the two-point correlation functions and the pairwise peculiar velocity dispersion for these galaxies.  A power-law fit to the real-space 2PCF $\\xi(r) =(r_0/r)^\\gamma$ gives an exponent $\\gamma=1.69$, and a amplitude $r_0=3.70 \\mpc$. We show that these results can be very well reproduced from mock samples constructed for the LCDM model with the CLW bias of $\\alpha = 0.25$. The bias needed for the PSCz galaxies is stronger than the LCRS galaxies, since IRAS galaxies tend to avoid high-density cluster regions. Mock smaples from the GIF simulation with an appropriate choice of galaxies give similar 2PCF.  The pairwise peculiar velocity dispersion measured from the PSCz has a much lower value, about $300 \\kms$ at $r_p = 1 \\mpc$, than the LCRS result at that separation of $570 \\pm 80 \\kms$.  All the simulation models which are consistent with the 2PCF of the PSCz predict significantly larger values for PVD, although the CLW bias reduces the PVD to within the $1 \\sigma$ limit of this value, at least near $3 \\mpc$ (if $\\Omega_0=0.2$).  As discussed in the last section, the decrease of $\\sigma_{12}(r_p)$ for $r_p \\ls 1\\mpc$ for the PSCz data is significant, and cannot be reproduced by the current simple models. After our paper was submitted to the journal, we noted that the steep decline of the PVD at the small speraration was also found in a recent work of the SDSS team on clustering of blue galaxies (Zehavi et al. 2002).  Thus, our results indicate that either the cosmogony model needs to be revised or the velocity bias is important for the velocity dispersion of the IRAS galaxies. We are investigating this issue with high-resolution SPH/N-body simulations."
    },
    "0201/astro-ph0201373_arXiv.txt": {
        "abstract": "We propose a general method for mapping the extinction in dense molecular clouds using 2MASS near-infrared data. The technique is based on the simultaneous utilization of star counts and colors. These two techniques provide independent estimations of the extinction and each method reacts differently to foreground star contamination and to star clustering. We take advantage of both methods to build a large scale extinction map ($2.5^\\circ \\times 2.5^\\circ$) of the North America-Pelican nebulae complex. With $K_s$ star counts and $H-K_s$ color analysis the visual extinction is mapped up to 35~mag. Regions with visual extinction greater than 20 mag account for less than 3\\% of the total mass of the cloud. Color is generally a better estimator for the extinction than star counts. Nine star clusters are identified in the area, seven of which were previously unknown. ",
        "introduction": "The knowledge of the extinction in molecular clouds is essential for a wide range of purposes. The extinction traces the dust distribution and thus provides an estimate of the column density of hydrogen that can be compared to the far-infrared or submillimeter emission from dust and to molecular emission lines. Fundamental properties are derived from these comparisons. \\citet{SAB+01} have found evidence of dust evolution in a translucent cloud in Taurus using extinction from $J$-2MASS star counts and submillimeter data from the balloon-borne experiment PRONAOS. \\citet{KAL+99} have compared the C$^{17}$O line with extinction derived from $H-K$ color and have found a depletion of CO for $A_V>10$ which indicates the interaction between dust an gas in high density regions. The dust distribution also gives information about the fragmentation processes which may be responsible for the shape of the initial mass function \\citep{PNJ97} through turbulent motions in the interstellar medium \\citep{PJN97}. All these results were obtained with near-infrared data which are better adapted than optical wavelengths for extinction studies. 2MASS provides accurate ($\\sim 3$\\%) $J$, $H$ and $K_s$ photometry for the whole sky. The $K_s$ band is about 10 times less absorbed than the $V$ band making possible to analyze the dust distribution for visual extinctions as large as 40~mag. Since the dust both absorbs and reddens the starlight, the analysis of the star density and star color variations give two independent estimates of the extinction. In this paper, we take advantage of these two techniques to measure the extinction and understand better the limitations of each method.  The first section of this paper describes the methods for mapping the dust distribution based on $K_s$ star counts and $H-K_s$ reddening. The comparison of the two maps leads to the identification of foreground stars and stellar clusters. In a second section, we apply our technique to the North America-Pelican complex which is located in the galactic plane at an intermediate distance ($\\sim 580$~pc). By analyzing the statistical and systematic uncertainties in the star count and in the color extinction maps we merge these two maps into a more reliable representation of the cloud. We also propose the identification of new stellar clusters in this complex field. In the appendix, we discuss the comparison between 2MASS data and a synthetic model for stellar populations. ",
        "conclusions": "We have proposed a method to map extinction with 2MASS data using simultaneously $K_s$ star counts and $H-K_s$ reddening. Finding a value of the extinction significantly larger from the count map than from the color map indicates significant foreground star contamination. Foreground stars can be removed individually in high extinction regions and statistically elsewhere. We stress the point that the foreground population dominates for high extinction and these regions would not be correctly analyzed if foreground stars are not carefully removed. On the other hand, a higher extinction value derived from colors than from counts indicates the presence of a star cluster. The adaptive method increases the angular resolution of the mapping around these objects. Color, especially $H-K_s$, is generally a better estimator of the extinction than star counts which have a higher statistical uncertainty.  Towards the NAN, we estimate the foreground star density to $1900\\pm 500$ stars~deg$^{-2}$ and the visual extinction is mapped up to 35~mag. Angular resolution of the extinction map depends on the local stellar density and is $1.2'$, $3.7'$ and $7'$ for visual extinctions of about 0, 20 and $>30$~mag, respectively. Nine clusters are identified, only two of which were already known and another two of which correspond to a diffuse excess of red stars close the center of the cloud, suggesting an {\\em evolved} young population or star forming activity in the whole core ($\\sim 5$~pc diameter).  Near-infrared wavelengths are required to investigate such highly obscured fields. With the help of the accurate 2MASS photometry, we have obtained results that supersede the previous work on the extinction in this area. The method developed here can be extended to other high extinction clouds with a significant fraction of foreground stars (i.e. located at more than 500~pc, in the galactic plane) that were not accessible before. The knowledge of the extinction toward the galactic plane is essential in order to constrain stellar population model parameters, such as the diffuse extinction and the giant star distribution, that are not well constrained at higher galactic latitudes (see appendix). The generalization of studies based on the comparison of the extinction (derived from star counts and color) with molecular lines and/or far--infrared and submillimeter emission is now possible with the 2MASS data releases."
    },
    "0201/astro-ph0201529_arXiv.txt": {
        "abstract": "Bolometric and 0.2-2 keV X-ray luminosities of the hot gas haloes of simulated disc galaxies have been calculated at redshift $z$=0. The TreeSPH simulations are fully cosmological and the sample of 44 disc galaxies span a range in characteristic circular speeds of $V_c$ = 130-325 km s$^{-1}$. The galaxies have been obtained in simulations with a considerable range of physical parameters, varying the baryonic fraction, the gas metallicity, the meta-galactic UV field, the cosmology, the dark matter type, and also the numerical resolution. The models are found to be in agreement with the (few) relevant X-ray observations available at present. The amount of hot gas in the haloes is also consistent with constraints from pulsar dispersion measures in the Milky Way. Forthcoming XMM and Chandra observations should enable much more stringent tests and provide constraints on the physical parameters. We find that simple cooling flow models over-predict X-ray luminosities by up to two orders of magnitude for high (but still realistic) cooling efficiencies relative to the models presented here. Our results display a clear trend that {\\it increasing} cooling efficiency leads to {\\it decreasing} X-ray luminosities at $z$=0. The reason is found to be that increased cooling efficiency leads to a decreased fraction of hot gas relative to total baryonic mass inside of the virial radius at present. At gas metal abundances of a third solar this hot gas fraction becomes as low as just a few percent. We also find that most of the X-ray emission comes from the inner parts ($r\\la 20$ kpc) of the hot galactic haloes. Finally, we find for realistic choices of the physical parameters that disc galaxy haloes possibly were {\\it more than one order of magnitude} brighter in soft X-ray emission at $z$$\\sim$1, than at present. ",
        "introduction": "In disc galaxy formation models infall of halo gas onto the disc due to cooling is a generic feature. However, the gas accretion rate and hot gas cooling history are at best uncertain in all models so far. It is thus not clear to which extent the gas cooling out from the galaxy's halo is replenishing that which is consumed by star formation in the disc. Such continuous gas infall is essential to explain the extended star formation histories of isolated spiral galaxies like the Milky-Way and the most likely explanation of the ``G-dwarf problem'' --- see, e.g., Rocha-Pinto \\& Marciel (1996) and Pagel (1997).  At the virial temperatures of disc galaxy haloes the dominant cooling mechanism is thermal bremsstrahlung plus atomic line emission. The emissivity, increasing strongly with halo gas density, is expected to peak fairly close to the disc and decrease outwards, and if the cooling rate is significant the X-ray flux may be visible well beyond the optical radius of a galaxy.  Recently, Benson \\etal (2000) compared ROSAT observations of three massive, nearby and highly inclined disc galaxies with predictions of simple cooling flow models of galaxy formation and evolution. They showed that these models predict about {\\emph{an order of magnitude}} more X-ray emission from the galaxy haloes (specifically from a 5-18 arcmin annulus around the galaxies) than observational detections and upper limits.  In this paper we present global X-ray properties of the haloes of a large, novel sample of model disc galaxies at redshift $z$=0. The galaxies result from physically realistic gravity/hydro simulations of disc galaxy formation and evolution in a cosmological context. We find that our model predictions of X-ray properties of disc galaxy haloes are consistent with observational detections and upper limits. Given the results of the theoretical models of Benson \\etal we list the most important reasons why simple cooling flow models over-predict the present day X-ray emission of disc galaxy haloes.  In section 2 we give a very short description of the disc galaxy simulations. In section 3 we briefly describe the X-ray halo emission calculations and in section 4 the results obtained. Section 5 constitutes the discussion and section 6 the conclusion. \\begin{figure} \\resizebox{\\hsize}{!}{{\\includegraphics{Tcut.ps}}} \\caption{The figure shows the temperature of the SPH gas particles in a typical disc galaxy from our simulations versus their x coordinate (one of the axes in the disc).  The ``cold'' (log$T<4.5$) gas which is primarily situated in the disc is removed from the catalogues since it does not contribute to the X-ray flux. } \\label{figT} \\end{figure} ",
        "conclusions": "We have presented X-ray properties of the hot gas haloes of disc galaxies derived from a large sample of physically realistic gravity/hydro simulations of galaxy formation and evolution.  The simulated galaxies follow an L$_{X,bol}$-$V_c$  relation with approximately the same slope as expected from simple cooling flow models ($L_{X,bol}\\propto V_c^5$), but shifted to lower $L_{X,bol}$, and with a significant scatter (approximately a 50\\% rms dispersion for a given choice of physical parameters).  The total bolometric X-ray luminosities of the disc galaxy haloes are up to two orders of magnitude less than predicted by simple cooling flow models. Hence, contrary to the order of magnitude discrepancy between simple cooling flow models and observations found by Benson \\etal (2000), our models are in agreement with  ROSAT observations of the three massive highly inclined spirals NGC 2841, NGC 4594 and NGC 5529. Furthermore, we find that our models are consistent with recent estimates of the diffuse 0.2-2 keV X-ray luminosity of the Milky Way, and also that the amount of hot gas in the haloes of our simulated, Milky Way sized disc galaxies is consistent with upper limits from pulsar dispersion measures toward the Magellanic Clouds and the globular cluster M53.  In contrast to what is predicted by simple cooling flow models, we find that {\\it increasing} the cooling efficiency  of the halo gas leads to a {\\it decrease} in the present day $L_X$. The reason for this is that increasing the cooling efficiency over the course of a simulation results in less hot gas in the halo at $z$=0 to cool (because the total amount of gas available at any given time is always limited to the gas inside of the virial radius). This in turn leads to lower  present day accretion rates and lower $L_{X,bol}$. The two most important physical parameters controlling the X-ray luminosities are the baryon fraction and the gas abundance: for a {\\it given} characteristic circular speed $V_c$, increasing the baryon fraction from $f_b$=0.05 to $f_b$=0.1 decreases $L_{X,bol}$ by a factor of about two, and similarly increasing the gas abundance from primordial to $Z=1/3~Z_{\\odot}$ results in a decrease in  $L_{X,bol}$ of a factor 3-4.  Concerning the spatial distribution of X-ray emission in the hot gas haloes we find that this is centrally concentrated: about 95\\% of the emission originates within the inner $r\\la$20 kpc.  For their $\\Lambda$CDM, $f_b$=0.1 and primordial gas abundance simulations Sommer-Larsen \\etal (2002) find present day mass accretion rates of 0.5-1 $M_{\\sun}$yr$^{-1}$ for Milky Way sized disc galaxies, compatible with observational limits (J. Silk, private communication). As the mass accretion rate effectively is proportional to $L_{X,bol}$ one would expect even lower accretion rates for the more realistic case of some level of metal enrichment of the hot gas. Furthermore, Sommer-Larsen \\etal (2002) find mass accretion rates which are an order of magnitude larger at redshift $z\\sim$1. Hence it is quite likely that disc galaxies were considerably X-ray brighter in the past.   Forthcoming XMM-Newton and Chandra observations of massive, nearby, edge-on disc galaxies will provide constraints on the present models, in particular the main physical parameters: the baryon fraction and the hot gas abundance. Future observations to redshifts $z\\ga$1 may be used  to constrain the models further."
    },
    "0201/astro-ph0201003_arXiv.txt": {
        "abstract": "We identify a subsample of the recently detected extrasolar planets that is minimally affected by the selection effects of the Doppler detection method. With a simple analysis we quantify trends in the surface density of this subsample in the period-$M\\sin (i)$ plane. A modest extrapolation of these trends puts Jupiter in the most densely occupied region of this parameter space, thus indicating that Jupiter is a typical massive planet rather than an outlier. Our analysis suggests that Jupiter is more typical than indicated by previous analyses. For example, instead of $M_{\\rm Jup}$ mass exoplanets being twice as common as  $2\\; M_{\\rm Jup}$ exoplanets, we find they are three times as common. ",
        "introduction": "The prevalence of infrared emission from accretion disks around young stars is consistent with the idea that such disks are ubiquitous. Their disappearance on a time scale of 50 to 100 million years suggests that the dust and gas accrete into planetesimals and eventually planets (Haisch \\etal 2001). Such observations support the widely accepted idea that planet formation is a common by-product of star formation (e.g. Beckwith  \\etal 2000). In the standard model of planet formation, Earth-like planets accrete near the host star from rocky debris depleted of volatile elements, while giant gaseous planets accrete in the ice zones ($\\gtrsim 4$ AU) around rocky cores (Lissauer 1995, Boss 1995). When the rocky cores in the ice zones reach a critical mass ($\\sim 10 \\: m_{\\rm Earth}$) runaway gaseous accretion (formation of jupiters) begins and continues until gaps form in the protoplanetary disk or the disk dissipates (Papaloizou \\& Terquem 1999, Habing \\etal 1999), leaving one or more Jupiter-like planets at $\\sim 4 - 10$ AU.  We cannot yet determine how generic the pattern described above is. However, formation of terrestrial planets is thought to be less problematic than the formation of Jupiter-like planets (Wetherill 1995). Gas in circumstellar disks around young stars is lost within a few million years and it is not obvious that the rocky cores necessary to accrete the gas into a Jupiter can form on that time scale (Zuckerman \\etal 1995). Thus, Jupiter-like planets may be rare. Planets may not form at all if erosion, rather than growth, occurs during collisions of planetesimals (Kortenkamp \\& Wetherill 2000). The present day asteroid belt may be an example of such non-growth. In addition, not all circumstellar disks produce an extant planetary system. Some fraction may spawn a transitory system only to be accreted by the central star along with the disk (Ward 1997). Also, observations of star-forming regions indicate that massive stars disrupt the protoplanetary disks around neighboring lower mass stars, aborting their efforts to produce planets (Henney \\& O'Dell 1999). Given these uncertainties, whether planetary systems like our Solar System are common around Sun-like stars and whether Jupiter-like planets are typical of such planetary systems, are important open questions.  The frequency of Jupiter-like planets may also have implications for the frequency of life in the Universe. A Jupiter-like planet shields inner planets from an otherwise much heavier bombardment by planetesimals, comets and asteroids during the first billion years after formation of the central star. Wetherill (1994) has estimated that Jupiter significantly reduced the frequency of sterilizing impacts on the early Earth during the important epoch $\\sim 4$ billion years ago when life originated on Earth. The removal of comet Shoemaker-Levy by Jupiter in 1994, is a more recent example of Jupiter's protective role.   To date (November, 2001), $74 $ giant planets ($M\\sin (i) < 13\\; M_{\\rm Jup}$) in close orbits ($< 4\\; AU$) around  66 nearby stars have been detected by measuring the Doppler reflex of the host star (Marcy \\etal 2001, Mayor \\etal 2001). Seven stars are host to multiple planets (six doubles, one triple system). Approximately 5\\% of the Sun-like stars surveyed possess such giant planets (Marcy \\& Butler 2000). The large masses, small orbits, high eccentricities and high host metallicities of these 74 exoplanets was not anticipated by theories of planet formation that were largely based on the assumption that planetary systems are ubiquitous and our Solar System is typical (Lissauer 1995).   Naef \\etal (2001) point out that none of the planetary companions detected so far resembles the giants of the Solar System. This observational fact however, is fully consistent with the idea that our Solar System is a typical planetary system. Fig. \\ref{fig:massperiodesp} shows explicitly that selection effects can easily explain the lack of detections of Jupiter-like planets. Exoplanets detected to date can not resemble the planets of our Solar System because the Doppler technique used to detect exoplanets has not been sensitive enough to detect Jupiter-like planets. If the Sun were a target star in one of the Doppler surveys, no planet would have been detected around it.  This situation is about to change. In the next few years Doppler planet searches will be making detections in the region of parameter space occupied by Jupiter. Thus \\begin{figure*}[!ht] \\centerline{\\psfig{figure=massperiodesp12.ps,height=13.0cm,width=15cm}} \\caption{Mass as a function of period for the 74 exoplanets detected to date. Regions where planets are ``Detected'', ``Being Detected'' and ``Not Detected'' by the Doppler surveys are shaded differently and represent the observational selection effects of the Doppler reflex technique (see Section \\ref{sec:partitions} for a description of these regions). The rectangle enclosing the grid of twelve boxes defines the subsample of $44$ planets less biased by selection effects. The numbers in the upper left of each box gives the number of planets in that box. The increasing numbers from left to right and from top to bottom are easily identified trends. In Figs. \\ref{fig:masshistogramfit} \\& \\ref{fig:periodhistogramfit} we quantify and extrapolate these trends into the lower mass bin and into the longer period bin which includes Jupiter. The ``+1'' and ``+5'' in the two boxes in the lower right refer to the undersampling corrections discussed in Section \\ref{sec:correction}. The seven exoplanetary {\\it systems} are connected by thin lines. Jupiter and Saturn are in the ``Not Detected'' region. The upper $x$ axis shows the distances and periods of the planets of our Solar System. The brown dwarf region is defined by $M\\sin (i)/M_{\\rm Jup} > 13$. The point size of the exoplanets is proportional to the eccentricity of the planetary orbits. } \\label{fig:massperiodesp} \\end{figure*} \\clearpage it is timely to use the current data to estimate how densely occupied that parameter space will be. The detected exoplanets may well be the observable  $5\\%$ tail of the main concentration of massive planets of which Jupiter is typical. The main goal of this paper is to correct or account for selection effects to the extent possible and then examine what the trends in the mass and period distributions indicate for the region of parameter space near Jupiter. Such an analysis is now possible because a statistically significant sample is starting to emerge from which we can determine meaningful distributions in mass, period as well as in eccentricity and metallicity. Our analysis helps answer the important question: How does our planetary system compare to other planetary systems?  In the next section we present our method for identifying a less biased subsample of exoplanets. In Section \\ref{sec:PMplane} we identify and extrapolate the trends in mass and period. In Section \\ref{sec:discuss} we discuss our analysis and compare our results to previous work. In Section \\ref{sec:summary} we summarize our results. ",
        "conclusions": "\\label{sec:discuss} \\subsection{Eccentricity}  A significant difference between the detected exoplanets and Jupiter, is the high orbital eccentricities of the exoplanets. The eccentricities of the planets of our Solar System were presumably constrained to small values ($e \\lsim 0.1$) by the migration through, and accretion of, essentially zero eccentricity disk material. A simple model that can explain the higher exoplanet eccentricities is that in higher metallicity systems, the higher abundance of refractory material in the protoplanetary disk may lead to the production of more planetary cores in the ice zone producing multiple Jupiters which gravitationally scatter off each other. Occasionnaly one will be scattered in closer to the central star and become Doppler-detectable (Weidenschilling \\& Marzari 1996). If that is the origin of the hot jupiters, then the detected exoplanets may be the high metallicity tail of a distribution in which our Solar System is typical, and as longer period giant planets are found they will have lower eccentricities, comparable to Jupiter's and Saturn's. Thus, if Jupiter is the norm rather than the exception, not only will we find more planets in the $P -M\\sin (i)$ parameter space near Jupiter as reported above, but also the eccentricities of the longer period exoplanets will be lower. The general distribution of the eccentricities of the exoplanets does not seem to reflect this, but exoplanets in planetary systems lend some support to the idea (see Fig. \\ref{fig:eccperiodesp} and caption).    \\clearpage \\begin{figure*}[!ht] \\centerline{\\psfig{figure=massperiodmore6.ps,height=13.0cm,width=15cm}} \\caption{The region of the $P - M\\sin (i)$ plane occupied by our Solar System compared to the region being sampled by Doppler surveys. Doppler surveys are on the verge of detecting Jupiter-like exoplanets. The lines, symbols and shading are the same as in Fig. \\ref{fig:massperiodesp}. We would like to know how planetary systems in general are distributed in this plane. Extrapolations of the trends quantified in this paper put Jupiter in the most densely occupied region of the $P - M\\sin (i)$ parameter space indicating that the detected exoplanets are the observable tail of the main concentration of massive planets that occupies the parameter space closer to Jupiter. If our Solar System is typical then the dispersion away from Jupiter into the Doppler-detectable region of this  plot may be largely due to the effect of high metallicity in producing gravitational scattering. Whether Jupiter is slightly more or less massive than the average most massive planet in a planetary system is difficult to determine. However, the strong correlation between the presence of Doppler-detectable exoplanets and high host metallicity (e.g. Lineweaver 2001) suggests that high metallicity systems preferentially produce massive Doppler-detectable exoplanets. This further suggests (since the Sun is more metal-rich than $\\sim 2/3$ of local solar analogues) that Jupiter may be slightly more massive than the average most massive planet of an average metallicity, but otherwise Sun-like, star. The dashed wedge-shaped contour represents the microlensing constraints discussed in the text. } \\label{fig:massperiodmore} \\end{figure*}  \\subsection{Fitting for $\\alpha$} \\label{sec:compare}  In this paper we have focused on the position of Jupiter relative to the exoplanets. This relative comparison does not require a conversion of exoplanet $M\\sin (i)$ values to $M$ values (Jorissen \\etal 2001, Zucker \\& Mazeh 2001a, 2001b and Tabachnik \\& Tremaine 2001). However, for this comparison Jupiter's position needs to be lowered and spread out a bit. Given a random distribution of orbital inclinations, the probability of $y =M \\sin (i)$, given $M$, is \\be P(y|M) = \\frac{y}{M^{2}\\sqrt{1-\\frac{y^{2}}{M^{2}}}}. \\ee With $M = M_{\\rm Jup}$, this probability is the curve placed outside the plotting area on the lower right of Fig. \\ref{fig:massperiodesp}. It represents the region of $M\\sin (i)$ that Jupiter-mass planets would fall in when observed at random orientations. The mean of this distribution is $\\pi/4 \\approx 0.79$ while the median is $0.87$ (in units of $M_{\\rm Jup}$). This lowers and spreads out in $M\\sin (i)$ the position of Jupiter but does not change the main results of the extrapolations done here.  The functional form $dN/d(M\\sin i) \\propto (M\\sin i)^{\\alpha}$ can be fit in various ways to various versions of the $M\\sin (i)$ histogram of exoplanets. When the histogram of all 74 exoplanets is fit, including the highly undersampled lowest $M\\sin (i)$ bin, the result is $\\alpha = -0.8$. This is reported in Marcy \\etal 2001 and we confirm this result. This value for $\\alpha$ is  close to the $\\approx -0.8 \\pm 0.2 $ found for very low mass stars (Bejar \\etal 2001). When the lowest exoplanet $M\\sin (i)$ bin is ignored because of known incompleteness we obtain $\\alpha = -1.1$. This is very similar to the  $\\alpha \\approx -1.1$ found in fits to the derived $M$ distribution (Zucker \\& Mazeh 2001a, Tabachnik \\& Tremaine 2001). The fit for $\\alpha$ seems to be more dependent on how the first bin is treated and how the sample for fitting is selected than on whether one fits to $M\\sin (i)$ or $M$.  Fitting to the $M\\sin (i)$ histogram of the less biased sample of 44 exoplanets, uncorrected for undersampling, yields $\\alpha = -1.3$. After correcting for undersampling as described in Section \\ref{sec:correction} we obtain our final result: $\\alpha = -1.5 \\pm 0.2$ (Fig. \\ref{fig:masshistogramesp}). This slope is steeper than the $\\alpha \\approx -1.0$ of previous analyses and indicates that instead of $M_{\\rm Jup}$ mass exoplanets being twice as common as  $2\\; M_{\\rm Jup}$ exoplanets, they are three times as common.   Despite the fact that massive planets are easier to detect, the mass distribution of detected planets is strongly peaked toward the lowest detectable masses. And despite the fact that short period planets are easier to detect, the period distribution is strongly peaked toward the longest detectable periods. To quantify these trends as accurately as possible, we have identified a less-biased subsample of exoplanets (Fig. \\ref{fig:massperiodesp}). Within this subsample, we have identified trends in $M\\sin (i)$ and period that are less biased than trends based on the full sample of exoplanets. Straightforward extrapolations of the trends quantified here, into the area of parameter space occupied by Jupiter, indicates that Jupiter lies in a region of parameter space densely occupied by exoplanets.  Our analysis indicates that 45 new planets will be detected in the parameter space discussed in the text. This estimate of 45 is a lower limit in the sense that  if a smooth curve, rather than our two straight boundaries, more accurately describes the selection effects in Fig \\ref{fig:massperiodesp}, larger corrections to the bin numbers would steepen the slopes in both Fig. \\ref{fig:masshistogramfit} \\&  \\ref{fig:periodhistogramfit}.  Despite the importance of the mass distribution and the trends in it, it is the trend in period that, when extrapolated, takes us to Jupiter and the parameter space occupied by Jupiter-like exoplanets (compare Figs. \\ref{fig:massperiodesp} \\& \\ref{fig:periodhistogramfit}). Long term slopes in the velocity data that have not yet been associated with planets are present in  a large fraction of the target stars surveyed with the Doppler technique (Butler, Mayor private communication). However, quantifying the percentage of host stars showing such residual trends is difficult and depends on instrumental noise, phase coverage and the signal to noise threshold used to decide whether there is, or is not, a long term trend.  Figure \\ref{fig:massperiodmore} shows that the Doppler technique has been able to sample a very specific high mass, short-period region of the $log P - log M\\sin (i)$ plane. Thus far, this sampled region does not overlap with the 10 times larger area of this plane occupied by the nine planets of our Solar System. Thus there is room in the $\\sim 95\\%$ of target stars with no Doppler-detected planets, to harbour planetary systems like our Solar System.  The trends in the exoplanets detected thus far do not rule out the hypothesis that our Solar System is typical. They support it. The extrapolations of the trends quantified here put Jupiter in the most densely occupied region of the $P - M\\sin (i)$ parameter space. Our analysis suggests that Jupiter is more typical than indicated by previous analyses -- instead of $M_{\\rm Jup}$ mass exoplanets being twice as common as $2\\; M_{\\rm Jup}$ planets we find they are three times as common. In addition long term trends in velocity, not yet identified with planets, are common. Both of these observations indicate that the detected exoplanets are the observable tail of the main concentration of massive planets of which Jupiter is likely to be a typical member rather than an outlier.  Null results from microlensing searches have been used to constrain the frequency of Jupiter-mass planets (Gaudi et al. 2002). These are plotted in Fig. \\ref{fig:massperiodmore}. Less than $33\\%$ of the lensing objects (presumed to be Galactic bulge M-dwarfs) have planetary companions within the dashed wedge-shaped area (the period scale, but not the AU scale, is applicable to this area). A conversion of the relative frequencies reported here to a fractional abundance in the wedge-shaped area yields the rough estimate that more than $\\sim 3$ percent of Doppler-surveyed Sun-like stars will be found to have companions with masses and periods in the wedge-shaped area. Thus our results are crudely consistent with current microlensing constraints. However, because of the difference in host mass, ($\\sim M_{Sun}$ for Doppler surveys and $\\sim 0.3 M_{Sun}$ for microlensing) it is not clear that such a direct comparison is meaningful. For example, if in the next few years Doppler and microlensing constraints appear to conflict, it may simply be that typical planetary masses scale with the mass of the host star, that is, Jupiter-mass planets at Jupiter-like orbital radii may be more common around $\\sim M_{Sun}$ stars than around $\\sim 0.3 M_{Sun}$ stars."
    },
    "0201/astro-ph0201079_arXiv.txt": {
        "abstract": "We present sensitive, high-resolution radio observations of the circumnuclear region of the barred spiral galaxy NGC6951. These observations reveal a ring of radio emission with many discrete components and a marginally resolved nuclear component. We compare the radio ring with observations at other wavelengths, and discuss the nature of the compact radio components. ",
        "introduction": "The dynamics and relationships between gas inflow processes, circumnuclear star formation and the presence of an active nucleus are important, since they might provide valuable leads towards understanding the transport of gaseous material towards the nuclear regions of active galaxies. An interesting class of objects for studying these aspects are the S\\'{e}rsic-Pastoriza or S-P galaxies (S\\'{e}rsic \\& Pastoriza 1965; S\\'{e}rsic 1973) which include galaxies with diffuse and amorphous nuclei in addition to the well-known hot-spot systems. S\\'{e}rsic \\& Pastoriza (1967) found that galaxies with these unusual nuclear morphologies occurred in barred or mixed-type galaxies. The S-P galaxies reflect a broad spectrum of properties with some of them also harbouring an active nucleus. The existence of both starburst and an active nucleus makes them particularly interesting for studying their kinematics, evolution and possible relationships between the two forms of activity. Circumnuclear star formation has been seen in many hot-spot galaxies, and there have been suggestions of a correlation between circumnuclear rings and nuclear activity (cf. Arsenault 1989).  Theoretical studies suggest that circumnuclear rings arise due to a bar-driven inflow of gas and dust to an inner Lindblad resonance (ILR) or between two ILRs (e.g. Combes \\& Gerin 1985; Athanassoula 1992; Byrd et al. 1994; Piner et al. 1995). The dense gas which accumulates in an ILR ring leads to a high star-formation rate either due to collisions of the molecular clouds (Combes \\& Gerin 1985), or gravitational collapse in the ring when the density reaches a critical value (Elmegreen 1994). To understand the formation and fuelling of the AGN, the gas must flow inwards from the ILR to sub-parsec scales, and this is at present not well understood (cf. Axon \\& Robinson 1996). Some simulations suggest small, steady inflow (Piner et al. 1995), while others suggest rapid inflow for brief periods (Wada \\& Habe 1992; Heller \\& Shlosman 1994).  In this paper we report high-resolution radio observations of the nuclear region in the barred, late-type grand-design galaxy NGC6951, which we had earlier observed as part of a survey of 47 S-P galaxies with the Very Large Array (VLA) at $\\lambda$20 and 6 cm (Saikia et al. 1994). The observations reported here are of higher resolution and better sensitivity, and clarify the radio structure of the circumnuclear region of this well-studied galaxy. New supernovae have been reported in this galaxy during 1999 and 2000 (Cao et al. 1999; Valentini et al. 2000).  \\begin{figure*}[t] \\vspace{0.2in} \\hbox{ \\hbox{ \\psfig{figure=fig1a.ps,width=3in} } \\hspace{0.8in} \\hbox{ \\hspace{-0.7in} \\psfig{figure=fig1b.ps,width=3in} } } \\caption{Radio images  of the nuclear region of NGC6951 at  8.4 GHz with an angular resolution of 2.57 $\\times$ 1.44 arcsec$^2$ along PA 0$^\\circ$ (left panel), and 0.86 $\\times$ 0.71 arcsec$^2$ along PA 113$^\\circ$ (right panel). The peak brightness and contour levels are indicated below each image. } \\end{figure*} ",
        "conclusions": "High-resolution radio observations of the circumnuclear region of NGC6951 reveal a ring of radio emission with many discrete components and a slightly resolved nuclear component. The discrete components have a median luminosity of 10$^{19}$ W Hz$^{-1}$ when observed with a linear resolution of $\\sim$90 pc, and are likely to consist of a mixture of both SNRs and \\hbox{H\\,{\\sc ii}} regions. The SN rate has been estimated to be $\\sim$0.07 yr$^{-1}$, which is reasonably consistent with those found in some of the archetypal starburst galaxies. The central component, which is resolved and has a steep radio spectrum, could be due to a small-scale jet from the active nucleus."
    },
    "0201/astro-ph0201286_arXiv.txt": {
        "abstract": "The microwave spectral energy distribution of the dusty, diffuse \\HII\\ region \\LPH\\ has been interpreted by others as tentative evidence for microwave emission from spinning dust grains. We present an alternative interpretation for that particular object; specifically, that an ultracompact \\HII\\ region embedded within the dust cloud would explain the available observations as well or better than spinning dust. Parameters for the size, surface brightness, and flux density of the putative ultracompact \\HII~region, derived from the microwave observations, are within known ranges. A possible candidate for such an ultracompact \\HII\\ region is \\irassrc, based upon its infrared colors. However, \\irassrc's infrared flux appears to be too small to be consistent with the microwave flux required for this alternative model to explain the observations. ",
        "introduction": "\\label{sec_intro}  Studies of the cosmic microwave background are affected by foreground emissions from the Milky Way and other sources. Principal mechanisms are thermal emission from warm dust, synchrotron from electrons gyrating in magnetic fields, and free-free (a.k.a. bremsstrahlung) emission from ionized plasma. An elusive forth foreground is attributed to spinning dust \\citep[]{draine98a} or perhaps magnetic dust \\citep[]{draine99}. A small but growing number of statistical comparisons of infrared tracers of dust and microwave observations of ``CMB quality'' have demonstrated repeatedly, but in each case at low significance, an excess of microwave emission, correlated with dust on the sky, and possibly due to spinning or magnetized dust. Finkbeiner {\\it et al.} (2002, Paper I) list such statistical comparisons and add to them two tentative detections of the spinning-dust mechanism in specific astronomical sources, \\LPH\\ and Lynds 1622, from a sample of 10 sources observed by them.  \\LPH\\ is the topic of this {\\it Letter}. It differs from the other nine dust clouds observed by Finkbeiner {\\it et al.} (2002) in that it is known to be a diffuse \\HII\\ region whereas the others are not ionized. Finkbeiner {\\it et al.} included \\LPH\\ in their sample despite their prejudice that free-free emission might overwhelm any emission from spinning dust. Alternatively, one might anticipate that ionized regions would be good places to find microwave emission from spinning dust, because in those regions, ion collisions with grains are expected to be the largest contributory factor by far in maintaining the spin of the grains \\citep[]{draine98b}.  Because of its special status as the sole \\HII\\ region in the list, and because \\LPH\\ provided the only very significant detection and was indeed much brighter than the spinning-dust theory would have predicted, we thought that a different explanation might exist. The alternative model for \\LPH\\ that we present in this {\\it Letter} is the superposition of one source with very large emission measure and very small angular scale with a low emission-measure, extended source. The combination of a small, optically thick source and a large optically thin source, both emitting by free-free, can create the rising microwave spectrum observed in Paper I. ",
        "conclusions": "\\label{sec_dis}  An optically thick ultracompact \\HII\\ region can provide a rising microwave spectrum with a spectral index $\\approx 2$ which is similar at these frequencies to the spinning-dust model's spectral index $\\approx 2.8$ (Ferrara \\& Dettmar 1994; Draine \\& Lazarian 1998b). For free-free emission, for $\\tau_{ff} \\ ^{>}_{\\sim} 1$ at $\\nu ^{<}_{\\sim} 10$ GHz, the EM must be $^{>}_{\\sim} 10^9$ cm$^{-6}$~pc.  While \\irassrc\\ is within a fraction of a 6\\arcmin\\ beam width of \\LPH, and has the infrared colors of an ultracompact \\HII\\ region, its \\IRAS\\ fluxes are too small unless its distance is improbably large. Also, the distance-independent ratio of the 100 $\\mu$m flux of \\irassrc\\ to our model's 2-cm flux is too small by more than an order of magnitude. The latter conflict is mitigated somewhat if the star is an early B-type star that has an atypically large ratio of ionizing to bolometric luminosity as has been observed in some cases (Cassinelli 1996) and has been predicted by modern stellar atmosphere models (Vacca, Garmany, \\& Shull 1996).  Paper I concluded that ``as long as the density is low, as in the LPH list, \\HII\\ regions may be optimal targets for future work for DASI, CBI, and \\MAP.'' It would be inappropriate to misinterpret this single case study of \\LPH\\ as contradictory to that recommendation. Instead, attention to its premise, that the density be low, is especially appropriate. The all-sky observations made by \\MAP\\ will allow astronomers to conveniently avoid any regions that might be unsatisfactory in that regard.  The hypothesis presented in this {\\it Letter} can be tested directly with radio interferometric imaging, which we plan to complete in early 2002. If that test shows not even one dense ionized knot in \\LPH, then this {\\it Letter}'s remaining value will be the analysis of Section \\ref{sec_modeling}, and the microwave emission from \\LPH\\ will still be a mystery."
    },
    "0201/astro-ph0201415_arXiv.txt": {
        "abstract": "We present the first visible wavelength images of the GG Tauri circumbinary ring, obtained with the Hubble Space Telescope's Wide Field and Planetary Camera 2.  Scattered light from the ring is detected in both V and I band images.  The images show that the ring is smooth, except for a small gap that could be a shadow caused by material between the stars and the ring.  The spokes seen extending from the stars to the ring in ground-based adaptive optics images are not seen in our data, which suggests that they may be image artifacts.  The nearside/farside surface brightness ratio is 6.9 in I band, consistent with forward scattering by small dust grains.  The azimuth of the peak ring surface brightness appears offset by 13$^{\\circ}$ from the azimuth closest to us, as seen in previous near-IR HST observations.  This may indicate that the ring is warped or somehow shadowed by the circumstellar disks.  The color of the ring is redder than the combined light from the stars as observed by HST, confirming previous measurements that indicate that circumstellar disks may introduce extinction of light illuminating the ring.  We detect a bright, compact arc of material 0\\farcs 3 from the secondary star at an azimuth opposite the primary.  It appears to be too large to be a circumstellar disk and is not at the expected location for dust trapped at a Lagrange point. ",
        "introduction": "GG Tauri (HBC 54; IRAS 04296+1725) is a young multiple star system located in the Taurus L1551 molecular cloud at a distance of 140 pc (Elias 1978).  It includes two binaries, GG Tau {\\it Aa/Ab}, with a separation of $\\sim$0\\farcs 25, and 10\\farcs 6 away, GG Tau {\\it Ba/Bb}, separated by 1\\farcs 5.  GG Tau {\\it Aa/Ab} (collectively hereafter just GG Tau) is especially interesting because it possesses a circumbinary disk that has been imaged in molecular line and dust continuum emission with millimeter interferometry (Dutrey, Guilloteau, \\& Simon 1994, hereafter DGS94; Guilloteau, Dutrey, \\& Simon 1999, hereafter GDS99).  Subsequently, images of the disk in scattered light were obtained in the near-infrared with ground-based adaptive optics (Roddier et al. 1996, hereafter R96) and the Hubble Space Telescope (HST) (Silber et al. 2000; McCabe \\& Ghez 2000).  These measurements show that the disk has a large inner hole, $\\sim$180 AU in radius, which has apparently been cleared by tidal interactions with the binary.  While the disk has been detected out to a radius of $\\sim$800 AU in millimeter line emission, virtually all of the scattered light and millimeter continuum emission is confined within an annulus of radius 180-260 AU.  For this reason, the disk is usually described as a ring.  The morphology is consistent with a circular ring inclined $\\sim$37$^{\\circ}$ from the plane of the sky.  The CO kinematics are in excellent agreement with Keplerian disk models (GDS99). Redshifted gas is found along the ring's western side, and blueshifted gas on its eastern.  Assuming that the disk and binary rotate in the same direction, this information and the observed motions of the stars indicate that the northern, brighter side of the ring must be the nearest to us (R96).  The adaptive optics (AO) images of R96, which were further analyzed by Close et al. (1998), showed a clumpy ring and indicated the presence of radial ``spokes'' extending inward from the ring toward the central binary.  Later AO images taken with Gemini (Potter et al. 2001) also contained some spokes.  These features resembled predictions for accretion streams from a circumbinary disk (Artymowicz \\& Lubow 1994).  However, both of the near-IR observations with HST (Silber et al. 2000; McCabe \\& Ghez 2000) show an essentially smooth ring, with no evidence for the spokes at the intensity levels indicated in the AO images.  Previous photometry of the individual binary components suggests that each may possess its own circumstellar disk.  Near-infrared excesses were noted in both stars by R96.  White et al. (1999) derived from HST spectra line-of-sight extinctions A$_V$=0.72 and 3.20 to GG Tau {\\it Aa} and {\\it Ab} respectively.  This large difference between two sources just 0\\farcs 25 (35 AU) apart (projected) shows that there must be extinction localized at one of the stars, perhaps from a small circumstellar disk.  Limits to the amount of circumstellar material are provided by high resolution 2.7 mm and 1.3 mm continuum maps by Looney, Mundy, \\& Welch (2000) and GDS99, respectively.  These show that no more than a few percent of the total flux from the system can originate in the circumstellar material of the individual binary components.  While any such disks must therefore be small and low mass in comparison to the outer circumbinary ring, their presence may have important effects on the ring's illumination. Indeed, R96 noted that the ring was redder than the combined light from the two stars and suggested that this indicated extinction along the line of sight between the central binary and the ring.  Wood, Crosas, \\& Ghez (1999) explored this scenario quantitatively, and through modeling of the scattered light images of R96 and $^{13}$CO J=1-0 line profiles of DGS94, found that extinction from circumstellar disk(s) appeared necessary to account for the brightness and color of scattered light from the ring. Despite the many lines of indirect evidence, no resolved images of circumstellar disks in the GG Tau binary have yet been obtained.  Based on its strong IR excess and millimeter continuum flux, GG Tau was included in the Wide Field and Planetary Camera 2 (WFPC2) Investigation Definition Team's HST observing program before the nature of the ring was confirmed by DGS94 and R96.  We present these results here, the first images of the GG Tau ring at visual wavelengths. ",
        "conclusions": "HST/WFPC2 V and I band images of GG Tau resolve the binary star and reveal the scattered light from the circumbinary dust ring that was previously detected in the near-IR.  The images confirm that the northern, nearer side of the ring is brightest, as is expected from forward scattering by dust grains.  The nearside/farside brightness ratio is 6.9 in the I band.  This effect has been reproduced using three-dimensional scattering models of the ring. The V-I$_c$ color of the ring relative to the binary colors indicates that some amount of extinction must be introduced by material between the stars, as first pointed out by R96.  The new images show a clump of reflecting material near the secondary star, also seen in the Silber et al. (2000) HST images (though apparently not noticed by them), that could possibly create shadows on, or at least introduce partial extinction to, the ring. Clear evidence for shadowing is provided by a narrow gap in the ring, which has also been seen in HST/NICMOS and adaptive optics images.  However, we cannot determine what may be causing this shadow, nor why it is only seen on one side if it is created by a circumstellar disk.  Other brightness asymmetries, such as the differences in the eastern and western ring radial profiles, and the offset of the ring peak brightness from the closest edge, might be explained by shadowing or possibly by warping of the ring.  The observations highlight the fact that the poorly understood interaction of light from the binary with ring material and possible circumstellar disks is the primary limitation to modelling the dust distribution as derived from scattered light images.  We find no evidence in our images for the ``spokes'' or ring clumpiness seen in the adaptive optics observations of R96.  We believe that these may be AO artifacts.  However, we do see some indication that the interior of the ring may not be completely clear, as there is a generally uniform distribution of light within the ring that is above the background level.  More observations are required to verify whether this is reflective material or simply a PSF subtraction residual.  Higher resolution imagery, which may provide details of the circumstellar disks, should substantially improve our understanding of GG Tau.  Ground-based adaptive optics imaging will help, but as demonstrated by the possible artifacts in the R96 images, any details will need to be confirmed by other systems.  In the near future, the Advanced Camera for Surveys (ACS) on HST will provide high resolution (0\\farcs 025 arcsec pixel$^{-1}$) visible- wavelength frames with a non-field-dependent PSF, which should improve PSF subtraction results compared to WFPC2.  GG Tau has been included in the ACS Science Team's program.  Further into the future, observations with ALMA (Atacama Large Millimeter Array) should provide high resolution (0\\farcs 01 - 0\\farcs 1) maps of disk emission, avoiding the complications introduced by shadowing and extinction in scattered light images."
    },
    "0201/astro-ph0201309_arXiv.txt": {
        "abstract": "\\targ\\ ($=$ RXJ1914.4+2457) shows pulsations in X-ray flux on a period of $9.5$ minutes, which have been ascribed to accretion onto a magnetic white dwarf, with the X-ray pulses seen as the accreting pole moves into and out of view. The X-ray flux drops to zero between pulses, and no other periods are seen, suggesting that \\targ\\ is a type of system known as a ``polar'' in which the white dwarf has a strong enough field to lock to the orbit of its companion.  If so, then \\targ\\ has the shortest orbital period known for any binary star. However, unlike other polars, \\targ\\ shows neither polarization nor line emission. In this paper we propose that \\targ\\ is the first example of a new type of X-ray emitting binary in which the mass transfer stream directly hits a non-magnetic white dwarf as a result of the very compact orbit. Our model naturally explains the X-ray and optical pulsations, as well as the absence of polarization and line emission. We show that direct impact will occur for plausible masses of the accreting star and its companion, e.g. $M_1 \\approx 0.5$, $M_2 \\approx 0.1\\,\\msun$. In our model \\targ\\ retains its status as the binary star with the shortest known orbital period, and is therefore a strong source of low-frequency gravitational waves. \\targ\\ is representative of an early phase of the evolution of the AM~CVn class of binary stars and will evolve into the normal disc-accretion phase on a timescale of $10^6$ to $10^7\\,\\yr$. The existence of \\targ\\ supports the double-degenerate route for the formation of AM~CVn stars. ",
        "introduction": "Accreting white dwarfs in binary stars fall into two groups according to the magnetic field of the white dwarf. The flow of material close to strongly magnetic white dwarfs is controlled by the magnetic field, and matter is channelled onto one or both magnetic poles. In this case, accretion energy is released as the material crashes into the white dwarf, emitting copious X-rays and optical cyclotron emission. The X-ray and cyclotron emission are modulated on the spin period of the white dwarf as the accreting poles rotate into and out of our view. Moreover, the cyclotron emission is circularly and linearly polarized. In the non-magnetic case, by contrast, the accreting material, having some initial angular momentum, forms a disc. The disc material accretes onto the white dwarf via an equatorial boundary layer where the kinetic energy of the gas is released. The resulting radiation is not significantly modulated, and there is no polarized cyclotron radiation. Pulsed X-ray emission from white dwarf binaries has therefore been regarded as a secure indication that the accretor is magnetic.  When X-ray pulsations on a period of $9.5\\,\\min$ were discovered from the star \\targ\\ (=RX J1914.4+2457), they were immediately interpreted in terms of an accreting, magnetic white dwarf spinning on the same period (Motch et~al. 1996).  Similar spin periods are commonly seen in the ``intermediate polar'' class of cataclysmic variable star in which a relatively weakly-magnetized white dwarf spins faster than the binary orbit. However, in these stars one also sees other periods, related to the orbital period and the ``beat'' period between the spin and orbital periods. In \\targ, however, both X-ray and optical data show just the one period of $9.5\\,\\min$. Moreover, the X-ray flux from \\targ\\ drops to zero between pulses, which is difficult to account for on an intermediate polar model. This led Cropper et~al. (1998) to suggest, instead, that \\targ\\ is a ``polar'', in which the white dwarf's magnetic field locks its spin to its companion star's orbit. If so, then \\targ\\ has an orbital period of only $9.5\\,\\min$, the shortest known for any binary system. Such a period implies that the donor is a helium-rich degenerate star, making \\targ\\ the first magnetic member of the AM~CVn stars. The AM~CVn stars are a select group of eight mass-transferring binary stars (other than \\targ), which have periods ranging from $17\\,\\min$ (AM~CVn, Nelemans, Steeghs \\& Groot 2001) to $65\\,\\min$ (CE~315, Ruiz et~al. 2001). As for \\targ, their short periods imply that the donor stars are hydrogen-deficient, and indeed no hydrogen appears in their spectra. They are thought to form from initially detached double white dwarf systems, from systems with helium-star donors or from mass transfer initiated when a $\\sim 1\\,\\msun$ donor starts to transfer mass to a white dwarf at the end of its core hydrogen burning (Nelemans et~al. 2001; Podsiadlowski, Han \\& Rappaport 2001). Double white dwarfs that fail to become AM~CVn stars are possible Type~Ia supernova progenitors.  A problem with the polar model is that \\targ\\ shows no optical polarization (Ramsay et~al. 2000), possible on the polar model only if the white dwarf has either a very strong or relatively weak field (while remaining synchronized). In an effort to explain this, Wu et~al. (2001) proposed a model in which the spin of the magnetic white is not synchronised with the orbit leading to dissipation of electric currents in the donor which produces unpolarized optical flux. However, the dissipation also leads to synchronization on a short timescale, which makes the chance of such a configuration low.  In this paper we present an alternative model for \\targ, in which the white dwarf need not be magnetic, but the X-rays will still be strongly modulated. \\targ\\ may thus be the first example of a new class of X-ray emitting binary star. We start by describing the observational characteristics of \\targ\\ that need explaining. ",
        "conclusions": "We have shown that plausible masses for the two components of \\targ\\ suggest that the mass transfer stream in this system strikes the accreting white dwarf directly. The resulting spot explains the X-ray pulses observed from this system, without the need for a magnetic white dwarf. This is consistent with the lack of polarization from the system, and with the complete disappearance of X-ray flux in between pulses. Our model has a lifetime of $10^6$ to $10^7\\,\\yr$ compared with the $10^3\\,\\yr$ of Wu et~al.'s unipolar inductor model, making it easier to reconcile with estimated formation rates for the AM~CVn systems, as long as a substantial fraction of these systems pass through a direct impact-phase, as predicted by Nelemans et~al. (2001) on the basis of a double white dwarf origin for these systems. \\targ\\ is the first example of a new class of stream-fed, non-magnetic white dwarf accretors."
    },
    "0201/astro-ph0201401_arXiv.txt": {
        "abstract": " ",
        "introduction": " ",
        "conclusions": "\\vspace{-4.0mm} Preliminary hydrodynamical and N-body simulations were undertaken with self-consistent star formation and gas heating/cooling. Our models successfully recover a pure gas Magellanic Stream, similar to that observed. Self-consistent treatments of star formation histories of the LMC and SMC are now underway.  This will rectify one of the remaining short comings of the models -- the near order-of-magnitude discrepancy between the mass of the simulated and observed Stream. These results represent the first self-consistent gas + N-body + star formation simulations of the Magellanic System.  \\vspace{-5.0mm}"
    },
    "0201/astro-ph0201292_arXiv.txt": {
        "abstract": "M16=NGC 6611, the Eagle Nebula, is a well studied region of star formation and the source of a widely recognized Hubble Space Telescope (HST) image. High spatial resolution infrared observations with the Near Infrared Camera and Multi-Object Spectrometer (NICMOS) on HST reveal the detailed morphology of two embedded star formation regions that are heavily obscured at optical wavelengths.  It is striking that only limited portions of the visually obscured areas are opaque at 2.2 microns.  Although the optical images imply substantial columns of material, the infrared images show only isolated clumps of dense gas and dust. Rather than being an active factory of star production, only a few regions are capable of sustaining current star formation. Most of the volume in the columns may be molecular gas and dust, protected by capstones of dense dust.  Two active regions of star formation are located at the tips of the optical northern and central large ``elephant trunk'' features shown in the WFPC2 images.  They are embedded in two capstones of infrared opaque material that contains and trails behind the sources. Although the presence of these sources was evident in previous observations at the same and longer wavelengths, the NICMOS images provide a high resolution picture of their morphology. Two bright stars appear at the tip of the southern column and may be the result of recent star formation at the top of that column. These observations suggest that the epoch of star formation in M16 may be near its endpoint. ",
        "introduction": "Dramatic optical pictures of M16 taken with WFPC2 on HST \\citep{hes96} show three large columns or elephant trunks of molecular gas and dust that are shaped by the stellar wind and radiation pressure from high luminosity young stars a few arc minutes away. The columns are designated I, II, and III going north to south.  The morphology of the columns clearly indicates that the initial high mass star formation in a molecular cloud region plays a large role in subsequent star formation by altering the density structure of the gas and dust in the surrounding region.  Large areas of low density are cleared of material while regions of higher density form elongated structures extending radially away from the cluster of newly formed stars.  The elongated trunk structures are repeated on a smaller scale with ``fingers'' and ``eggs'' of material throughout the larger structure.  It has generally been assumed that new star formation is occurring inside the trunk and egg structures and that gravitational attraction of the protostellar condensations is retaining the material against the action of the radiation pressure and stellar wind from the nearby young stars.  Earlier infrared observations by \\citet{hil93} indeed revealed significant previous star formation that was not easily visible by optical means, although they limited their quantitative analysis mainly to objects with detectable visible flux. They determined that there are stars up to 80 M$_{\\sun}$ in the visible OB association and that there are intermediate mass stars still above the main sequence, assuming their distance of 2 kpc which we will adopt in our analysis. The embedded sources discussed here are clearly visible in their image, but were not remarked upon individually as they were not detected in the optical image.  They also do not appear in their table of embedded infrared sources.  The non-stellar morphology of the sources is not clearly apparent in their published image, although there is nebulosity near the source locations.  There is also evidence for both sources (which we label M16ES-1 and M16ES-2 for M16 Embedded Source 1 and 2) in the WFPC 2 F547M images.  M16ES-2,  visible as a diffuse reflection nebulosity, was first described by \\citet{hes97}.  Although at lower resolution and much brighter limiting magnitude, M16ES-1 and 2 are clearly visible in the near infrared image of M16 by \\citet{mcc97}, however, they are not commented on in the text.  In this black and white rendering of a portion of a much larger 3 color (J,H,K) image the transparency of large sections of the columns is clearly apparent.  \\citet{mcc97} comments that at the limiting magnitude of the image only 10\\% of the eggs labeled by \\citet{hes96} appear to contain forming stars or chance coincidences with background sources.  Work done on the optical emission line images of the region \\citep{hes96} demonstrated the influence of photoevaporation on the morphology of the gas and dust.  In the absence of any other effects, photoevaporation will eventually destroy the region leaving only those stars that managed to form before the destruction.  The observations discussed in this paper are part of a program to study the detailed morphology of the embedded star formation in the trunks and to study the effects of embedded star formation on the subsequent evolution of the region.  A detailed analysis of the sources in the region is in preparation \\citep{smth02,hes02}. ",
        "conclusions": "There are at least two currently active areas of star formation in the columns of M16.  Both are near the tops of the columns which face the cluster of luminous O stars.  Although the optical images show significant extinction throughout the columns, the infrared images show that the density of dust and molecular gas is quite low, except at the location of the embedded sources and at isolated areas along the columns to the southeast. Far infrared images show another strong emission area to the southeast, beyond the extent of our images. It appears that this region of M16 is in its last stage of star formation and the dissipation of the columns may soon follow as the material in the few dense clumps is exhausted.  This explains the relatively low rate of star formation commented on by \\citet{pil98}.  An alternative is that further star formation will be triggered in the columns as the photocompressed region moves down the length of each column.  The dust and molecular gas in the columns away from the dust structures has a high enough density to be opaque to optical emission but not enough to produce high amounts of current star formation.  CO emission along the columns indicates that the amount of dust is sufficient to shield some molecular gas from dissociation.  The dense gas and dust at the locations of the two embedded sources act as capstones, shadowing much of the column from the ionizing radiation of the O star cluster. The O star cluster, however, has a significant angular extent, and ionizing radiation does fall on the surface of the columns.  The detailed interaction with the surface is described by \\citet{hes96}.  The lack of ionizing radiation from M16ES-1 indicates that it is either a cluster of relatively low mass stars or an object or objects that have not yet evolved to the Zero Age Main Sequence. The line emission in atomic hydrogen is due to photo-ionization and photo-excitation by radiation from the O star cluster. The morphology of the dust near the embedded sources suggests that they have a role in determining the local structure.  The elementary calculations of Sec.~\\ref{sec-loc} indicate that the sources have sufficient strength to affect their surroundings, but only in regions well shielded from the radiation from the OB association.  The restriction of the dust regions to a relatively narrow cone that points in the direction of the O star cluster shows that the cluster radiation plays the primary role in the dust and gas dynamics at the site of the embedded sources.  High resolution camera 2 images show two bright stars of very similar magnitude and color within the concave depression in the tip of column III.  These may be a chance superposition of local background stars on this location.  The stars, however, may have formed in the tip of the column.  If so the stars may have reduced the density in their surroundings enough to allow the photoevaporation from the O star cluster to be much more efficient in this area, forming the current concave depression in the column tip.  The three columns would then show a sequence of star formation stages from embedded in column I, emerging in column II, and fully emerged in column III."
    },
    "0201/astro-ph0201547_arXiv.txt": {
        "abstract": "The likelihood ratio test (LRT) and the related $F$ test,\\footnote{The $F$ test for an additional term in a model, as defined in \\cite{bevington:69} on pp. 208-209, is the ratio $$F_\\chi=\\frac{\\chi^2(m)-\\chi^2(m+1)}{\\chi^2(m)/(N-m-1)}= \\Delta\\chi^2/\\chi^2_\\nu,$$ \\noindent where $\\chi^2(m)$ and $\\chi^2(m+1)$ are the values of $\\chi^2$ statistic resulting from fitting $m$ and $m+1$ free parameters respectively and $\\chi^2_\\nu$, in Bevington's (1969) notation, stands for a $\\chi^2$ random variable with $\\nu$ degrees of freedom divided by the number of degrees of freedom $\\nu$. In the remainder of this paper we use $\\chi^2_\\nu$ to denote a $\\chi^2$ random variable with $\\nu$ degrees of freedom, as this notation is more standard.} popularized in astrophysics by \\citet{eadie:etal}, \\citet{bevington:69}, \\citet{lampton:margon:bowyer:76}, \\citet{cash:79}, and \\citet{avni:etal} do not (even asymptotically) adhere to their nominal $\\chi^2$ and $F$ distributions in many statistical tests common in astrophysics, thereby casting many marginal line or source detections and non-detections into doubt. Although the above references illustrate the many legitimate uses of these statistics, in some important cases it can be impossible to compute the correct false positive rate. For example, it has become common practice to use the LRT or the $F$ test for detecting a line in a spectral model or a source above background despite the lack of certain required regularity conditions.  (These applications were not originally suggested by \\citet{cash:79} or by \\citet{bevington:69}). In these and other settings that involve testing a hypothesis that is on the boundary of the parameter space, {\\it contrary to common practice, the nominal $\\chi^2$ distribution for the LRT or the $F$ distribution for the $F$ test should not be used}.  In this paper, we characterize an important class of problems where the LRT and the $F$ test fail and illustrate this non-standard behavior.  We briefly sketch several possible acceptable alternatives, focusing on Bayesian posterior predictive probability-values.  We present this method in some detail, as it is a simple, robust, and intuitive approach.  This alternative method is illustrated using the gamma-ray burst of May 8, 1997 (GRB 970508) to investigate the presence of an Fe~K emission line during the initial phase of the observation.  There are many legitimate uses of the LRT and the $F$ test in astrophysics, and even when these tests are inappropriate, there remain several statistical alternatives (e.g., judicious use of error bars and Bayes factors).  Nevertheless, there are numerous cases of the inappropriate use of the LRT and similar tests in the literature, bringing substantive scientific results into question. ",
        "introduction": "Distinguishing a faint spectral line or a new source from a chance fluctuation in data, especially with low photon counts, is a challenging statistical task.  As described in Section~2, these are but two examples in a class of problems that can be characterized in statistical terms as a test for the presence of a component in a finite mixture distribution. It is common practice to address such tests with a likelihood ratio test (LRT) statistic or the related $F$ statistic and to appeal to the nominal asymptotic distributions or {\\it reference distribution}\\footnote{As detailed below, the reference distribution is used to calibrate a test statistic. When choosing between two models, we assume the simpler or more parsimonious model holds and look for evidence that this assumption is faulty.  Such evidence is calibrated via the {\\it reference distribution}, the known distribution of the test statistic under the simple model. If the observed test statistic (e.g., LRT or $F$ test) is extreme according to the reference distribution (e.g. $\\chi^{2}_{1} > 10.83)$, the simple model is rejected in favor of the more complex model.} of these statistics \\cite[etc.]{mura:etal:88, feni:etal:88,   % yosh:etal:92, palm:etal:94, band:etal:95, band:etal:96, band:etal:97, freeman:etal:99, piro:etal:99}.  (See \\citet{band:etal:97} for a discussion of the close relationship between the LRT and the $F$ test.)  The underlying assumption is that in some asymptotic limit the statistic being used to describe the data is distributed in an understandable way, and hence useful bounds may be placed on the estimated parameters.  Unfortunately, the standard asymptotic theory does not always apply to goodness-of-fit tests of this nature even with a large sample size or high counts per bin. Thus, use of these statistics may be misleading.  For example, searches for cyclotron scattering and atomic lines in $\\gamma$-ray bursts based on such uncalibrated statistics may be unreliable.  In nested\\footnote{E.g., the allowed parameter values of one model must be a subset of those of the the other model.} significance tests such as the LRT or $F$ test, we wish to choose between two models, where one model (the null model) is a simple or more parsimonious version of the other model (the alternative model).  We seek evidence that the null model does not suffice to explain the data, by showing the observed data are very unlikely under the assumption that the null model is correct.  In particular, we define a test statistic (e.g., the LRT statistic, $F$ statistic, or $\\Delta \\chi^2$ values) with a distribution that is known at least approximately assuming the null model is correct.  We then compute the test statistic for our data and compare the result to the known null distribution.  If the test statistic is extreme (e.g., large) we conclude the null model is very unlikely to produce such a data set and choose the alternative model. A test statistic {\\it without a known reference distribution} is without standard justification and is of little direct use; we can neither calibrate an observed value nor compute the false positive rate.  Such is the case for the LRT statistic and the $F$ statistic for detecting a spectral emission line, an absorption feature, or {\\it added} model component\\footnote{We assume that when testing for the presence of an emission line, model parameters are constrained so the line intensity is greater than zero; similar assumptions are made for absorption features and added model components.}.  Because this use is outside the bounds of the standard mathematical theory, the reference distribution is generally {\\it uncalibrated: unknown and unpredictable}. This problem is fundamental, i.e., intrinsic to the definition of the LRT and $F$ test.  It is not due to small sample size, low counts per bin, or faint signal--to--noise ratio.  It persists even when Gauss--Normal statistics hold and $\\chi^2$ fitting is appropriate.\\footnote{As pointed out by the referee, \\citet{whea:etal:95} showed that least squares or $\\chi^2$ fitting can sometimes be equivalent to maximum likelihood fitting even when Poisson statistics apply.  However, their method is not universally applicable since it presumes that the weight matrix is independent of (or only weakly dependent on) the Poisson means (see their Equations 12, 19, and 20), whereas in the Poisson case the weights are the reciprocal of the Poisson means---if the weights are known, there is nothing to estimate.}  Several authors \\citep{mattox:etal:96, deni:wald:99,freeman:etal:99} have recognized that the null distribution of the LRT and $F$ statistics may vary from the nominal tabulated values \\citep[e.g., the tables given in][for the $F$ test]{bevington:69}.  Nonetheless, the inappropriate use of the $F$ test in the astrophysics literature is endemic.    As a rule of thumb, there are two important conditions that must be satisfied for the proper use of the LRT and $F$ statistics.  First, {\\it the two models that are being compared must be nested}. Second, {\\it the null values of the additional parameters may not be on the boundary of the set of possible parameter values}.  The second condition is violated when testing for an emission line because the line flux must be non-negative and the null value of this parameter is zero, which is the boundary of the non-negative numbers.  Because of the first condition, it is, for example, inappropriate to use the $F$ test to compare a power law with a black body model or to compare a power law model with Raymond Smith thermal plasma model. This issue is discussed in \\citet{freeman:etal:99} and is not the primary focus of this paper.  Instead, we focus on the second condition that disallows, for example, testing for an emission line, an absorption feature, or other {\\it added} spectral components (e.g., a power law, Compton reflection component, black body component, thermal component, etc.) or testing for a quasi-periodic oscillation in timing data, an added Gaussian in light curves, or an {\\it added} image feature (e.g., an added point source).  We emphasize that there are many legitimate uses of the LRT and $F$ test, e.g., testing for a broken power law; comparing the variances of two samples; determining if a spectrum is the same or different in two parts of an image, at two difference times, or in two observations; or deciding whether to allow non-solar abundances. Generally, returning to the two rule-of-thumb conditions should guide one as to when the $F$ test and LRT are appropriate.  (We note, however, that the reference distributions of these tests are only reliable with a sufficiently large data set, even when the two conditions are met.)  In the remainder of the paper we explain why these standard tests fail and offer alternatives that work.  In Section~2, we look at the class of models known as finite mixture models (which allow for multiple model components).  We discuss the yet unresolved question of determining the number of components in the mixture and show that testing for a spectral line, a new source, or other {\\it added} model component are special cases of this problem.  The LRT and the $F$ test have often been proposed as simple solutions in these cases. The LRT is specifically discussed in Section~3, and is shown to be invalid (i.e., uncalibrated) {\\it in this setting} since as we discussed above its basic criteria are not met. In Section~4, we discuss a number of possible alternatives to the LRT and $F$ statistics including Bayes factors, the Bayesian Information Criterion, and posterior predictive p-values.\\footnote{A probability-value or p-value is the probability of observing a value of the test statistic (such as $\\chi^2$) as extreme or more extreme than the value actually observed given that the null model holds (e.g. $\\chi^2_{30} \\geq 2.0$) Small p-values are taken as evidence against the null model; i.e., p-values are used to calibrate tests.  Posterior predictive p-values are a Bayesian analogue; see Section~4.2.}  Complete recipes for all alternatives to the LRT and $F$ test are beyond the scope of this paper.  As discussed in Section~4, we focus on posterior predictive p-values because they are conceptually and computationally simple and closely mimic the nested significance tests described above; see the high--redshift quasar example in Sections~3 and 4.  With this machinery in place, we investigate a typical example in Section~5; we investigate whether the data support the presence of an Fe K emission line during the initial phase of the GRB 970508.  Here neither the LRT nor the $F$ test are appropriate.  On the basis of our analysis, the model with a spectral line is clearly preferable. In Section~6, we conclude with several cautions regarding the inappropriate use of statistical methods which are exemplified by the misuse of the LRT.  Throughout the paper we use the LRT for a spectral line as an example, but the discussion and results apply equally to related tests (e.g., the $F$ test) and to other finite mixture models (e.g. how many sources are detected above background; \\citet{avni:etal}) or even more generally to testing any null models on the boundary of the parameter space. ",
        "conclusions": ""
    },
    "0201/astro-ph0201221_arXiv.txt": {
        "abstract": "We present constraints on the stellar--mass distribution of distant galaxies.  These stellar mass estimates derive from fitting population--synthesis models to the galaxies' observed multi-band spectrophotometry.  We discuss the complex uncertainties (both statistical and systematic) that are inherent to this method, and offer future prospects to improve the constraints.  Typical uncertainties for galaxies at $z\\sim 2.5$ are $\\delta(\\log \\mathcal{M}) \\sim 0.3$~dex (statistical), and factors of $\\mathrel{\\hbox{\\rlap{\\hbox{\\lower4pt\\hbox{$\\sim$}}}\\hbox{$>$}}} 3$ (systematic). By applying this method to a catalog of NICMOS--selected galaxies in the Hubble Deep Field North, we generally find a lack of high--redshift galaxies ($z\\mathrel{\\hbox{\\rlap{\\hbox{\\lower4pt\\hbox{$\\sim$}}}\\hbox{$>$}}} 2$) with masses comparable to those of present--day ``$L^\\ast$'' galaxies.  At $z\\mathrel{\\hbox{\\rlap{\\hbox{\\lower4pt\\hbox{$\\sim$}}}\\hbox{$<$}}} 1.8$, galaxies with $L^\\ast$--sized masses do emerge, but with a number--density below that at the present epoch.  Thus, it seems massive, present--day galaxies were not fully assembled by $z\\sim 2.5$, and that further star formation and/or merging are required to assemble them from these high--redshift  progenitors.  Future progress on this subject will greatly benefit from upcoming surveys such as those planned with \\emph{HST}/ACS and \\emph{SIRTF}. ",
        "introduction": "With current observations and those of the near future, we are able to observe distant galaxies ($z\\mathrel{\\hbox{\\rlap{\\hbox{\\lower4pt\\hbox{$\\sim$}}}\\hbox{$>$}}} 2$) in their primeval stages, i.e., at an era when they are vigorously assembling their stellar content. However, no conclusive picture has yet emerged to describe how these high--redshift galaxies fit into the ancestral history of the present--day galaxy population.  By measuring the stellar--mass distribution (which contains a complete historical record of star formation) for galaxies as a function of redshift, one can directly probe the global, mass--assembly history.  This provides a stringent test for cosmological models that recount how high--redshift galaxies evolve into the present--day galaxy population.  However, a galaxy's stellar mass is \\emph{not} a directly measurable quantity: it must be inferred from models of the galaxy's mass--to--light ratios and the observed multi-band photometry.  In this contribution, we discuss the method used to obtain stellar--mass estimates of distant galaxies and some the underlying caveats inherent in the process.  We then present results from applying this method to a NICMOS--selected sample of galaxies in the \\emph{Hubble Deep Field North} (HDF--N). ",
        "conclusions": "\\begin{figure} \\begin{center} \\leavevmode \\includegraphics[width=0.75\\textwidth]{papovichF4.eps} \\end{center} \\caption{Distribution of stellar mass for galaxies in the HDF--N as a function of co-moving volume. The stellar--mass estimates assume models with solar metallicity, Salpeter IMF, and single--component star--formation histories.  Solid symbols denote galaxies with spectroscopically confirmed redshifts, and open symbols those galaxies where only photometric redshifts are available.  The horizontal dashed line indicates the characteristic stellar mass of a present--day $L^\\ast$ galaxy~\\cite{col01}.  The solid curve traces the ``mass--limit'' for a maximally old galaxy formed as a burst at $z\\sim \\infty$ with passive evolution, and normalized to the flux of the NICMOS detection limit ($H_\\mathrm{AB} \\approx 26.5$). } \\label{fig:hdfmass} \\end{figure}  Although at present the stellar--mass estimates for high--redshift galaxies' have significant uncertainties, these constraints are interesting nevertheless.  For LBGs with ``$L^\\ast$'' UV luminosities \\cite{ste99}, we infer stellar mass estimates of $\\sim 10^{10}\\;\\mathcal{M}_\\odot$ or $\\sim 1/10$th that of a present--day $L^\\ast$ galaxy \\cite{col01}.  Extending this analysis to all galaxies in the NICMOS HDF--N catalog allows a comparison between the LBG population and galaxies down to more modest redshifts ($z\\mathrel{\\hbox{\\rlap{\\hbox{\\lower4pt\\hbox{$\\sim$}}}\\hbox{$>$}}} 0.5$). In fig.~\\ref{fig:hdfmass}, we show the distribution of galaxy stellar mass in the HDF--N as a function of co-moving volume.  Here, all stellar masses assume solar metallicity, a Salpeter IMF, and use only the single--component star--formation histories.  As such, they are nominally strict lower limits.  Also shown in the figure is a fiducial curve denoting the minimal detectable stellar mass of a maximally old galaxy as a function of redshift and the NICMOS detection limit.  Old galaxies would be detectable with masses \\emph{above} this curve. This, however, does not limit the minimal detectable masses of galaxies with lower mass--to--light ratios.  \\begin{figure} \\begin{center} \\leavevmode \\includegraphics[width=0.75\\textwidth]{papovichF5.ps} \\end{center} \\caption{Color--magnitude diagram for the HDF--N galaxies with $1.9 \\leq z \\leq 3.5$.  Solid symbols denote galaxies with spectroscopically confirmed redshifts, and the open symbols those galaxies with only photometric redshifts.  The solid curves denotes the evolution of a $10^{10}\\;\\mathcal{M}_\\odot$ galaxy at $z=2.7$ formed in a $\\delta$--function star--formation history. Note that the ``J''--dropout, HDF--N J123656.3+621322 \\cite{dic00}, is the only candidate for an old, red galaxy in the HDF--N in this redshift range.} \\label{fig:vhcmd} \\end{figure}  There are several interesting implications from fig.~\\ref{fig:hdfmass}.  Firstly, the HDF--N exhibits a lack of of $\\mathcal{M} \\mathrel{\\hbox{\\rlap{\\hbox{\\lower4pt\\hbox{$\\sim$}}}\\hbox{$>$}}} \\mathcal{M}^\\ast(z=0)$ galaxies at $z\\mathrel{\\hbox{\\rlap{\\hbox{\\lower4pt\\hbox{$\\sim$}}}\\hbox{$>$}}} 2$. Such galaxies should be detected (if present) in the deep NICMOS data, even to $z \\sim 3$ (beyond which the NICMOS $H$ band shifts below the 4000~\\AA/Balmer break and the stellar mass estimates are less secure). However, as shown in fig.~\\ref{fig:vhcmd}, there are few (if any) galaxies in this redshift range with $V_{606}-H_{160}$ colors indicative of a galaxy dominated by old stellar populations.  Thus, it is unlikely that we are missing them if they were present (however, see recent results from the HDF--S, e.g., Labb\\'e et al., this volume).  It is a possibility that we have underestimated their stellar masses due to the uncertainties described above.   Secondly, by $z\\mathrel{\\hbox{\\rlap{\\hbox{\\lower4pt\\hbox{$\\sim$}}}\\hbox{$<$}}} 1.8$, the upper envelope of stellar mass in the HDF--N increases to include massive, ``$L^\\ast$''--sized galaxies. Thus, it seems that the stellar populations of the progenitors to the massive galaxy population do not appear to be fully assembled in $z\\mathrel{\\hbox{\\rlap{\\hbox{\\lower4pt\\hbox{$\\sim$}}}\\hbox{$>$}}} 2$ progenitors.  This in turn suggests that more star--formation or merging (or both) are required for $z\\mathrel{\\hbox{\\rlap{\\hbox{\\lower4pt\\hbox{$\\sim$}}}\\hbox{$<$}}} 2$ to construct the large--galaxy population observed at $z\\mathrel{\\hbox{\\rlap{\\hbox{\\lower4pt\\hbox{$\\sim$}}}\\hbox{$<$}}} 1$ and at the present--epoch.  We wish to thank the conference organizers for arranging such a stimulating meeting in a beautiful setting. Support for this work was provided by NASA through grant GO--07817.01-96A."
    },
    "0201/astro-ph0201017_arXiv.txt": {
        "abstract": "We present a direct measurement of the stellar velocity dispersion of the early-type lens galaxy D in the system MG2016+112 ($z=1.004$), determined from a spectrum obtained with the {\\sl Echelle Spectrograph and Imager} (ESI) on the W.M. Keck--II Telescope, as part of the {\\sl Lenses Structure and Dynamics (LSD) Survey}. We find a velocity dispersion of $\\sigma_{\\rm ap}=304\\pm27$ km\\,s$^{-1}$ inside an effective circular aperture with a radius of $0.65''$, corresponding to a central velocity dispersion of $\\sigma=328\\pm32$ km\\,s$^{-1}$. From a {\\sl Hubble Space Telescope} F160W--band image, we measure the effective radius and effective surface brightness in order to determine the offset of the lens galaxy with respect to the local Fundamental Plane. The offset corresponds to an evolution of the rest-frame effective mass-to-light ratio of $\\Delta \\log(M/L_B)=-0.62\\pm0.08$ from $z=0$ to $z=1.004$. By interpreting colors and offset of the FP with two independent stellar population synthesis models, we obtain a single-burst equivalent age of 2.8$\\pm$0.8~Gyr (i.e.~$z_{\\rm f}$$>$1.9) and a supersolar metallicity of $\\log [Z/Z_\\odot]$=0.25$\\pm$0.25. The lens galaxy is therefore a massive elliptical dominated by an old and metal rich stellar population at $z>1$. The excellent agreement of the stellar velocity dispersion with that predicted from recent lens models confirms that the angular separation of the multiple images of the background QSO is predominantly due to the lens galaxy, and not to a massive ``dark cluster'', in agreement with recent weak lensing and X--ray observations. However, the significant overdensity of galaxies in the field might indicate that this system is a proto-cluster, in formation around galaxy D, responsible for the $\\sim$10\\% external shear inferred from the strong lens models. ",
        "introduction": "In the cold dark matter (CDM) cosmological scenario, structures in the Universe form through hierarchical merging of smaller structures (White \\& Rees 1978; Blumenthal et al.\\ 1984; Davis et al.\\ 1985). Within this general framework, early-type galaxies (E/S0) in the cores of rich clusters form at high redshifts ($z>2$), corresponding with the first dark-matter overdensity peaks, as opposed to a later formation epoch for field E/S0 (Kauffmann 1996).  Clusters subsequently form around these seeds by accretion of smaller-mass galaxies, with significant structural and dynamical evolution occurring at more recent cosmological times. Recent observations have shown that massive cluster E/S0 were already assembled at $z>1$ and subsequently evolved passively through ageing of their stellar populations (e.g. van Dokkum et al.\\ 1998; Stanford, Eisenhardt \\& Dickinson 1998). Similarly, field E/S0 galaxies seem not to evolve dramatically between $z=1$ and $z=0$, both in number (Schade et al.\\ 1999, Im et al.\\ 2002) and structural properties (van Dokkum et al.\\ 2001; Treu et al.\\ 1999, 2001b, hereafter T01b), although secondary episodes of star formation might be frequent at $z<1$ (Menanteau et al.\\ 2001; Treu et al.\\ 2002, hereafter T02).  However, most observational results on the evolution of E/S0 concern the evolution of their stellar populations and little is known about the evolution of their internal structure.  To comprehensively quantify the luminosity, color {\\sl and} structural evolution of the stellar mass and of the dark-matter halo of E/S0 as function of redshift, we are conducting an observational program with the Echelle Spectrograph and Imager (ESI) on the W.M.~Keck--II Telescope: the {\\sl Lenses Structure and Dynamics (LSD) Survey}. Aim of the {\\sl LSD survey} is to measure the internal kinematics of a dozen gravitational-lens galaxies up to $z=1$, allowing a powerful combination of dynamical and lensing constraints on their mass distribution. The LSD Survey\\footnote{see also http://www.its.caltech.edu/\\~{}koopmans/LSD or http://www.astro.caltech.edu/\\~{}tt/LSD}, its main goals and observing strategy will be described in detail elsewhere (Treu \\& Koopmans 2002, in preparation).  Here we present the first result of the {\\sl LSD Survey}, a measurement of the stellar velocity dispersion of the lens galaxy in the system MG2016+112 at $z=1.004$. A summary of relevant prior observations and a new lens model can be found in Koopmans et al.\\ (2002) and references therein. The primary lens galaxy (D) in MG2016+112 is the highest spectroscopically-confirmed redshift lens galaxy known to date. The suggestion that this lens was embedded in a massive ``dark cluster'', based on ASCA X-Ray observations (Hattori et al. 1997), was recently shown to be incorrect by high-resolution Chandra observations, showing no evidence for hot X-ray gas (Chartas et al. 2001). On the other hand, deep optical studies that show an overdensity of at least six bright E/S0 with the same redshift as the lens galaxy (Benitez et al. 1999; Soucail et al. 2001; Clowe et al. 2001).  The absence of a significant weak-lensing signal (Clowe et al. 2001) and the absence of X-ray emission, however, shows that these galaxies are not associated with a massive evolved cluster.  In the following, we assume for definiteness that the Hubble constant, the matter density, and the cosmological constant are H$_0=65$~km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_{\\rm m}=0.3$, and $\\Omega_{\\Lambda}=0.7$, respectively. ",
        "conclusions": "We have presented a direct measurement of the stellar velocity dispersion of the lens galaxy in the highest spectroscopically-confirmed redshift lens system MG2016+112 ($z=1.004$), as part of the {\\sl Lenses Structure and Dynamics (LSD) Survey}. The effective radius and surface brightness are determined from HST images, to compare the properties of the lens galaxy with those of E/S0 in the local Universe. The colors and evolution of M/L$_B$ of its stellar population, as derived by comparison to the local FP, indicate an old ($\\sim$3\\,Gyr) metal-rich stellar population, consistent with a passively evolved single stellar population formed at $z_{\\rm f}>1.9$.  The large stellar velocity dispersion and metallicity of the lens galaxy is comparable to the most massive present-day cluster galaxies (Ferreras et al. 1999; Fisher et al.\\ 1995). In addition, the F814W-band surface brightness profile at large radii deviates slightly from the $R^{1/4}$ profile, possibly indicating an outer envelope, as typically found for brightest cluster galaxies and cD galaxies (Graham et al.\\ 1996). This bluer envelope and the [OII] emission that we detect from galaxy D (see also Soucail et al. 2001) could be due to the accretion of a younger stellar populations from smaller gas-rich galaxies, such as the closest companion to galaxy D for which we detect [OII] in emission (Fig.2).  \\begin{inlinefigure} \\begin{center} \\resizebox{\\textwidth}{!}{\\includegraphics{f4.eps}} \\end{center} \\figcaption{Likelihood contours of the age and metalicity of the stellar population of MG2016+112, galaxy D. The inner and outer contours indicate the 68\\% and 95\\% probability. Panels (a) and (b) are for population synthesis models from Bruzual \\& Charlot (1993) and Fioc \\& Rocca-Volmerange (1997), respectively. \\label{plotfour}} \\end{inlinefigure}   This scenario is consistent with the gravitational-lens system being embedded in a galaxy overdensity (i.e. a proto-cluster; see also Benitez et al.\\ 1999), which is not yet relaxed to a more centrally concentrated system and still shows evidence for ongoing merging and accretion. In addition, it would explain the absence of a significant weak-lensing signal (Clowe et al.\\ 2001), the relatively small lensed-image separation, as well as constraints from recent Chandra X-ray observations (Chartas et al.\\ 2001). In particular, the spread of the velocities of the galaxies in the overdensity (Soucail et al.\\ 2001), if they are assumed to be virialized, implies a mass three times larger than the 3--$\\sigma$ upper limit inferred from the non-detection of extended X-ray emission with Chandra (Chartas et al. 2001).  We therefore think that the assumption of virialization is {\\sl not} correct and that these galaxies form a (non-virialized) proto-cluster. Three galaxies in the field are spatially clustered at $z=0.97$ and might be falling towards the overdensity with high velocity (Soucail et al. 2001), consistent with a cluster in formation.  ~\\\\ \\begin{inlinetable} \\centering \\begin{tabular}{lr} \\hline \\hline Redshift (D)          & 1.004$\\pm$0.001 \\\\ $F160W$ (mag)      & 18.24$\\pm$0.02 \\\\ $F814W$--$F160W$ (mag)     & 3.3$\\pm$0.1 \\\\ SB$_{e,F160W}$  (mag/arcsec$^2$)& 17.64$\\pm$0.40 \\\\ $R_{e,F814W}$ (arcsec) & 0.65$\\pm$0.10 \\\\ $R_{e,F160W}$ (arcsec)& 0.31$\\pm$0.06 \\\\ $\\sigma_*(< 0.65'')$  (km\\,s$^{-1}$) &\t304$\\pm$27 \\\\ $b/a$=$(1-e)$ & 0.75$\\pm$0.10 \\\\ Major axis P.A. ($^\\circ$) & $121\\pm2$\\\\ \\hline $\\sigma$  (km\\,s$^{-1}$) &\t328$\\pm$32 \\\\ $B-I$ (mag)    & 1.98$\\pm$0.15 \\\\ $M_{B}$ (mag)  & $-$22.53$\\pm$0.10 \\\\ $M_{I}$ (mag)  & $-$24.51$\\pm$0.04 \\\\ SB$_{e,B}$ (mag/arcsec$^2$) & 18.12$\\pm$0.50 \\\\ SB$_{e,I}$  (mag/arcsec$^2$)& 16.14$\\pm$0.40 \\\\ \\hline \\hline \\end{tabular} \\end{inlinetable}  \\noindent{Table~1.--- Observed spectro-photometric quantities of galaxy D in MG2016+112.  The second part of the table lists rest-frame quantities, derived from the observed quantities (see text).}    {"
    },
    "0201/astro-ph0201367_arXiv.txt": {
        "abstract": "Up to now, only a very small number of dwarf novae have been studied during their outburst state ($\\sim$30 per cent in the Northern hemisphere). In this paper we present the first comprehensive atlas of outburst spectra of dwarf novae. We study possible correlations between the emission and absorption lines seen in the spectra and some fundamental parameters of the binaries. We find that out of the 48 spectra presented, 12 systems apart from IP~Peg show strong He{\\sc ii} in emission: SS~Aur, HL~CMa, TU~Crt, EM~Cyg, SS~Cyg, EX~Dra, U~Gem, HX~Peg, GK~Per, KT~Per, V893~Sco, IY~UMa, and 7 others less prominently: FO~And, V542~Cyg, BI~Ori, TY~Psc, VZ~Pyx, ER~UMa, and SS~UMi. We conclude that these systems are good targets for finding spiral structure in their accretion discs during outburst if models of Smak (2001) and Ogilvie (2001) are correct. This is confirmed by the fact that hints of spiral asymmetries have already been found in the discs of SS~Cyg, EX~Dra and U~Gem. ",
        "introduction": "Dwarf novae are a type of cataclysmic variable that undergo outbursts during which they increase in brightness by 2--8 mags. The outburst recurrence time varies widely from system to system and can be as short as a few days (ER~UMa systems) and as long as tens of years (WZ~Sge systems). The most likely mechanism for dwarf nova outbursts is a thermal instability within an accretion disc (Osaki 1974). There are still many unknowns in the disc instability theory. For example, although we know that the viscosity in the disc plays a fundamental role in driving the outburst, we do not know what its origin is. In 1986 Sawada, Matsuda \\& Hachisu suggested that the disc viscosity could be due to shocks in the disc tidally excited by the donor star. In 1991, Balbus \\& Hawley discovered a strong magneto-hydrodynamical instability of accretion discs which can indeed produce considerable effective viscosity. But it is not yet certain which, if either, of these is dominant.  There are several extensive studies on dwarf novae in the literature (Warner 1995 and references therein) but they concentrate mainly on their quiescent states.  Only $\\sim$30 per cent of known Northern hemisphere dwarf novae have published outburst spectra and the number for Southern hemisphere systems is even smaller. The main aim of this paper is to fill this gap by presenting spectra of 48 dwarf novae during their outburst state. We carry out comparisons between the spectra of the different dwarf novae and search for possible correlations between the features seen in the spectra and fundamental parameters of the systems such as their inclination, the mass of their components, their orbital periods and the outburst phases at which they were observed. Comparisons between the spectra are possible because the data were taken with comparable setups during a long term programme using service time.  Another reason for compiling this atlas is to look for spectra similar to that of IP~Peg, U~Gem and WZ~Sge, systems that have been confirmed to show spiral structure in their discs during outburst. By confirmed we mean that spiral structure has been detected in more than one set of data, either taken during a different outburst or during the same outburst at a different time (IP~Peg: Steeghs et al. 1997, Harlaftis et al. 1999, Morales-Rueda, Marsh \\&\\ Billington 2000; U~Gem: Groot 2001, Steeghs private communication; WZ~Sge: Kuulkers et al. 2001, Baba, Sadakane \\&\\ Norimoto 2001). The principal feature of their spectra, compared with those of other dwarf novae, is that they show strong lines in emission instead of absorption. Of the $\\sim$30 per cent of dwarf novae with outburst spectra published, only half of these show any emission, and even then it is often in the form of emission lines buried in absorption troughs. There could easily be systems with strong emission lines in outburst that have yet to be identified. Of most interest is the presence of \\heii\\ in emission as this line is very important from the irradiation point of view. If these asymmetric structures seen in the discs of IP~Peg, U~Gem and WZ~Sge are the result of tidally thickened sectors of the disc being irradiated by the white dwarf, boundary layer and/or inner disc during the outburst (as suggested by Smak 2001 and Ogilvie 2001) we expect them to be seen most clearly in \\heii. ",
        "conclusions": "We present the first comprehensive sample of dwarf nova spectra during outburst. We study the possible correlations between the lines present in the spectra and fundamental parameters of the systems.  We find that there is no correlation between the strength of the lines and the masses of the components of the system or the periods of the binaries. There seems to be some correlation between the strength of the lines and the inclination of the system as represented by the FWHM of the lines. We do not find any correlation between the FWHM or the EW of the lines and either the time of the outburst at which the dwarf novae were observed or the outburst recurrence time measured for each system.  Of the 48 dwarf novae outburst spectra presented we find that 20 systems, 13 of them very clearly, and the other 7 possibly, show \\heii\\ in emission. If models presented by Smak (2001) and Ogilvie (2001) are correct, these are good candidates to show spiral structure in their accretion discs during outburst. These results are even more encouraging because four of these candidates: SS~Cyg, EX~Dra, U~Gem and IP~Peg, have been seen to show hints of spiral structure in their discs already (Steeghs et al. 1996; Joergens, Spruit \\& Rutten 2000, Steeghs, Harlaftis \\& Horne 1997, Groot 2001). Based on recent outburst spectra of U~Gem (Steeghs, private communication) we suggest that in fact the \\heii\\ emission is either the result of irradiation or produced by the spiral structures themselves. If its origin is irradiation we expect the \\heii\\ intensity to change with disc thickness. If the origin of \\heii\\ is the structures themselves, we should see them in all dwarf novae that show \\heii\\ emission.  Time resolved spectroscopy of the other candidates found in our sample will allow us to confirm or dismiss the presence of spiral asymmetries in their accretion discs.  \\subsection*{Acknowledgements}  LM-R was supported by a PPARC post-doctoral research grant. The Isaac Newton telescope is operated on the island of La Palma by the Isaac Newton Group in the Spanish Observatorio del Roque de los Muchachos of the Instituto de Astrof\\'{\\i}sica de Canarias. The authors would like to thank the Variable Star Network, VSNET, and the Variable Star Observers League in Japan, VSOLJ, for making their data available as most of the light curves presented in this paper were produced using their observations. In this research, we have used, and acknowledge with thanks, data from the AAVSO International Database, based on observations submitted to the AAVSO by variable star observers worldwide. The reduction and analysis of the data was carried out on the Southampton node of the STARLINK network. We also wish to thank the referee and D. Steeghs for useful comments and L. Gonz{\\' a}lez Hern{\\' a}ndez for computer support.  \\begin{table*} \\vbox to220mm{\\vfil Landscape table to go here \\caption{} \\vfil} \\end{table*}"
    },
    "0201/astro-ph0201151_arXiv.txt": {
        "abstract": "We calculate the spectrum of ultra high energy cosmic rays produced by the decay of a superheavy dark matter population clustered in the galactic halo. To perform this calculation we start with fragmentation functions measured at LEP and evolve them to the cosmic ray energy scale using the QCD DGLAP equations. We consider Standard Model evolution and supersymmetric evolution. We also take into account many--body final states in the decay of the dark matter particles. ",
        "introduction": "Ultra High Energy Cosmic Rays (UHECR) are microscopic particles --protons, photons or perhaps more exotic objects-- with a macroscopic energy, about 50 Joules per particle in the extreme of the present observed spectrum. They strike the upper layers of the Earth atmosphere at a rate of about one event per century and per kilometer square. Over one hundred of them have been detected so far by kilometer scale detectors and many more are expected to be seen by forthcoming observatories.  Explaining how these particles get this huge energy is a real challenge for our present understanding of the cosmos. One UHECR primary gets somehow ten million times the energy that a proton will gain at the future Large Hadron Collider (LHC) at CERN. Nature has always been at ease beating the achievements of mankind. One can envisage two broad classes of mechanisms by which nature can produce UHECRs, the so-called bottom up class models and the top-down class models.  In bottom-up models charged particles are accelerated by magnetic fields in large astrophysical sites. In top-down models particles are not accelerated but are created at birth with the huge energy typical of UHECRs.  Here we will concentrate on a particular top-down model~\\cite{chor,bkv97,bs98}. We will assume the UHECRs are produced by the decay of a population of superheavy dark matter particles with a lifetime longer than the age of the Universe $\\tau > 10^{10}$~yr and with mass $\\sim 100$ Joule$/c^2$. Theoretical motivation for these particles can be found in~\\cite{crypton}.  A superheavy dark matter population created at some stage of the early universe~\\cite{gravprod,reheatprod} will gravitationally cluster in the galactic halo. Since the length scale of the halo is around 100 kpc, UHECRs produced in the galactic halo will not have time to interact with the CMB before they reach the Earth. The absence of GZK cut-off is a genuine prediction of models where UHECRs are produced by the decay of a superheavy dark matter particle clustered in the galactic halo.  Since the halo shape is close to spherical one expects a quasi isotropic distribution of events from a halo superheavy dark matter population, which is compatible with experiments. At present it is not possible to make any strong claim about the observed angular distribution of UHECR events because of low statistics. Hopefully, future observatories like Pierre Auger will gather a large enough sample of events to settle down this issue~\\cite{efs01}.  Finally there is the question of the UHECR composition. The cosmic ray observatories cannot measure the composition of the primary flux (whether they are photons, protons or heavy nuclei) in an event by event basis. Composition can only be determined in a statistical way. Present analyses tend to favour protons as the main component in the primary flux~\\cite{ahvwz00}. In top-down models the main component in the production site, in our case the the galactic halo, is the photon component (actually, neutrinos dominate over photons in a large range of the spectrum but their probability to be detected is too small with present or past detectors). However, on their way to the Earth, photons will interact with the low frequency radio background in the Galaxy and their total flux may be substantially diminished so that on the Earth the total photon flux may be comparable or smaller than the baryon flux. Whether this is indeed possible, subject to the EGRET bound on the low energy $\\gamma$-rays which will be created by the electromagnetic cascading, will be discussed elsewhere.  Summing up, there are three main tests to falsify or support the hypothesis of UHECRs produced by the decay of dark matter particles clustered in the galactic halo. The first one is the energy distribution of events or spectrum. The second one is the expected angular distribution of events in the sky. The third one is the composition of the primary flux. Here we will focus on the first test and, partially, on the third one. We will briefly show how quantum chromodynamics (QCD) can be used to calculate the spectrum and composition (without photon galactic processing) of the expected UHECR flux. For further details see~\\cite{SarkarToldra,Toldra}. QCD was also used to calculate the spectra of UHECRs in~\\cite{Rubin,FodorKatz,BerezKachel01}. ",
        "conclusions": "We have calculated the UHECR spectra of baryons, photons and neutrinos expected from the decay of superheavy dark matter particles of mass $M_X\\sim 10^{12}$~GeV which are clustered in the galactic halo. The spectrum for every primary particle is given by FFs at the scale $M_X$. We have calculated these FFs using QCD DGLAP equations. We have considered Standard Model equations and SUSY equations. We have also taken into account the possibility of many-body decay in the decay of $X$. The shape of the fragmentation spectrum (of either baryons or photons) fits rather well the new component of ultra-high energy cosmic rays extending beyond the GZK energy."
    },
    "0201/astro-ph0201420_arXiv.txt": {
        "abstract": "X-ray reflection spectra from photoionized accretion discs in active galaxies are presented for a wide range of illumination conditions. The energy, equivalent width (EW) and flux of the \\fe\\ line are shown to depend strongly on the ratio of illuminating flux to disc flux, \\fxfd, the photon index of the irradiating power-law, $\\Gamma$, and the incidence angle of the radiation, $i$. When \\fxfd$ \\leq 2$ a neutral \\fe\\ line is prominent for all but the largest values of $\\Gamma$. At higher illuminating fluxes a He-like \\fe\\ line at 6.7~\\kev\\ dominates the line complex. With a high-energy cutoff of 100~\\kev, the thermal ionization instability seems to suppress the ionized \\fe\\ line when $\\Gamma \\leq 1.6$. The \\fe\\ line flux correlates with \\fxfd, but the dependence weakens as iron becomes fully ionized. The EW is roughly constant when \\fxfd\\ is low and a neutral line dominates, but then declines as the line progresses through higher ionization stages. There is a strong positive correlation between the \\fe\\ EW and $\\Gamma$ when the line energy is at 6.7~\\kev, and a slight negative one when it is at 6.4~\\kev. This is a potential observational diagnostic of the ionization state of the disc. Observations of the broad \\fe\\ line which take into account any narrow component would be able to test these predictions. Ionized \\fe\\ lines at 6.7~\\kev\\ are predicted to be common in a simple magnetic flare geometry. A model which includes multiple ionization gradients on the disc is postulated to reconcile the results with observations. ",
        "introduction": "\\label{sect:intro} The discovery of the iron K$\\alpha$ line and Compton reflection in the X-ray spectra of accreting black holes was an important step in the understanding of the central engine \\citep*{mu93}. These features were predicted to result from the reprocessing of X-rays by an optically thick and cold medium \\citep*{lw88,gr88,geo91,ma91}, quite likely the accretion disc itself. This was spectacularly confirmed by the \\textit{ASCA} observation of a broad \\fe\\ line in the X-ray spectrum of the bright Seyfert~1 galaxy \\mcg\\ \\citep{tan95}. The profile of this line was well fit by a model of line emission from relativistically moving material within 10 Schwarszchild radii of a supermassive black hole \\citep*{fab89}. Alternative explanations for such a broad line suffer from physical inconsistencies and/or fine-tuning of model parameters \\citep*{fab95,reyw00}. Suddenly a potentially powerful probe of accretion and black hole physics was observationally accessible \\citep*{fab00}. Many other active galactic nuclei (AGN) were subsequently observed by \\textit{ASCA} in search of a broad \\fe\\ line, and, although most of the detections were of far less quality than the one of \\mcg, the mean line profile of a sample of Seyferts seemed to be broadened in a similar manner \\citep*{n97a,yaq02}.  Since AGN themselves have quite variable X-ray continua, it was anticipated that the \\fe\\ line should also change over an observation. The largest dataset has come from long observations of \\mcg\\ where the line is seen to sometimes drastically change energy \\citep*{iwa96,iwa99}, and its flux varies independently of the continuum \\citep{ve01}. This last property, also seen in other objects \\citep*{wan99,ngm99,chi00,wan01,wea01}, is contrary to the predictions of the simple reflection scenario, which predicts the line variations should track the continuum. The fact that this is not observed suggests a more complicated and dynamic pattern of illumination on the disc \\citep{rey00} such as in the model of irradiation by magnetic flares in a patchy corona \\citep*{gal79,haa91,haa94,mf01}. These flares may be rotating \\citep{rus00}, outflowing \\citep*{bel99,mbp01} and/or temporally connected \\citep*{pf99,mf01}, and so they will clearly have an impact on the observed \\fe\\ line due to the changing pattern of radiation on the surface of the disc.  Many models of X-ray reflection from AGN-like accretion discs have been published \\citep*[e.g.][]{ros93,zyk94,mz95,ros99}, but they have not progressed to the same dynamic level as the models of accretion disc coronae. The best current models compute the reflection spectrum from a photoionized layer on the surface of an accretion disc in hydrostatic equilibrium \\citep*{nkk00,brf01,rdc01}. If discs are anything like what standard theory predicts, then these calculations should be able to make specific predictions about reflection spectra \\citep[e.g.][]{nay00,nk01}.  Here, we compute such spectra from an AGN accretion disc over a wide range of illumination conditions, and examine the behaviour of the energy, equivalent width (EW) and flux of the \\fe\\ line. Our approach differs from that of \\citet{nk01} or \\citet{zyk01} by concentrating on the observationally accessible changes to the \\fe\\ line.  The parameters are chosen so that conditions appropriate for both flaring and quiescent regions of the disc are explored. Comparison of the model predictions with time-averaged data (to avoid any non-equilibrium effects) may allow constraints to be placed on how discs are irradiated.  The paper is structured as follows. Section~\\ref{sect:comp} describes the model, the assumptions, and the range of parameters that are treated. The results of the calculations and a discussion of the evolution of the \\fe\\ line profile, flux and EW are given in Section~\\ref{sect:res}. A discussion of the results is given in Section~\\ref{sect:discuss} before Section~\\ref{sect:concl} summarizes the primary conclusions. ",
        "conclusions": "\\label{sect:concl} The broad \\fe\\ line which (most likely) results from reflection from an accretion disc is the most important spectral signature in the X-ray study of AGN. The model calculations presented in this paper have shown that the behaviour and properties of the Fe line from an irradiated disc in equilibrium depend on how it is illuminated. Specifically, when the incident flux is larger than the disc flux, we predict that a He-like \\fe\\ line at 6.7~\\kev\\ will be a common feature in the reflection spectrum. The thermal ionization instability is able to change this and produce a neutral line at 6.4~\\kev, but only when $\\Gamma \\leq 1.6$ (for a cutoff energy of 100~\\kev) and the incident flux is not too high.  The \\fe\\ EW will remain roughly constant as \\fxfd\\ is increased until the atmosphere becomes highly ionized and then it will decrease to small values. The rate of decline is lower for higher values of $\\Gamma$. When the line complex is dominated by He-like Fe, the EW will strongly increase with $\\Gamma$, but will slightly decline with the photon-index if it is neutral. This effect could be used as an observational diagnostic to determine the ionization state of the line, independent of the line centroid.  Long observations of Seyferts \\& quasars by \\textit{XMM-Newton} would be required to test the predictions of the models. However, any narrow component to the line profile must be taken into account before a comparison could be made. The results of such a study would be important, as it is clear that the ionization state of the disc depends greatly on how it is irradiated which would greatly constrain coronal models.  The prediction that ionized lines are common can be reconciled with their observational rarity by employing multiple ionization gradients on the accretion disc, which would be a natural result of a magnetically active and patchy corona. The challenge is now extended to modellers to compute observationally testable predictions for such complex scenarios."
    },
    "0201/astro-ph0201051.txt": {
        "abstract": " ",
        "introduction": "Proximity Focussing Ring Imaging counters (PFRICH) are based on a very simple geometrical configuration. The counter principle consists of a simple thin solid or liquid radiator, separated from the photodetector plane by a gap allowing photon rings associated to \\cer cones to expand and reach a suitable radius before they are detected (see \\cite{FRICH,YPS} for a general overview of RICH counters). \\par The price to pay for this architectural simplicity is a modest velocity resolution of the counter with respect to the best achievable performances \\cite{HYPO}. This type of configuration is suitable for counter designs requiriring a large geometrical acceptance \\{detection area\\}$\\otimes$\\{angular range\\}, for which the use of focussing devices is severely limited \\cite{FRICH} or even impracticable, provided the required velocity resolution is not too high. The limiting resolution of these counters is set by the chromatic dispersion of the radiator material. In practice, the thickness of the radiator used as well as the spatial resolution of the photodetector array are also limiting factors to the counter resolution. The issue has been extensively discussed in a previous report on a simulation study of the counter \\cite{SIMU} which complements the present work. Some of the results of this study will be repeated here for the reader's convenience. \\par The AMS project consists of a particle spectrometer scheduled to be installed aboard the International Space Station (ISS) by the year 2004 for a 3 to 5 years campaign of measurements, with a broad physics program \\cite{AMS,BE10}. The spectrometer will include a RICH counter among its instruments. The purpose of this counter is to achieve particle identification with the resolution performances shown to be realistic in the simulation study for mass and charge measurements. These are: \\\\ a) A one $amu$ (atomic mass unit) mass separation for light nuclei over a broad momentum range extending from about 1~GeV/c per nucleon, up to around 13~GeV/c per nucleon  at best, for mass numbers A$\\approx$20. This could be obtained by combining two radiators as shown in \\cite{SIMU}. \\\\ b) A one charge unit separation for nuclei up to $Z\\approx$~25 at best, for charge measurements, over the full momentum range of the spectrometer, i.e., from threshold up to above 1~TeV per nucleon. This latter performance depends critically on the electronics and PMT gain stability and calibration. \\par The full geometrical acceptance (${\\cal S}\\cdot\\Omega$, ${\\cal S}$ area, $\\Omega$ angular acceptance) of the spectrometer will be of the order of $\\approx$0.5~$m^2\\cdot sr$ for the RICH $\\otimes$ TOF $\\otimes$ TRACKER combination of detectors. The overall spectrometer dimensions are restricted by rigid constraints on the payload envelope which must fit inside the space shuttle bay. These requirements were pointing to a PFRICH type solution because of its simplicity, although alternative more ambitious options could have been taken. \\par The counter described here was a study prototype of first generation, built to perform an end-to-end test of the technique, from implementation of each component involved, up to velocity measurement, including the (first generation) prototype of front-end electronics, and event reconstruction algorithm. The purpose was to get through all the steps of the experimental procedure and to uncover all unexpected difficulties in order to finally reach the stage of the final counter design with a proven technique. The main points were: 1) investigating velocity and charge resolution capabilities of the counter over the full range of acceptance, in particular for large particle trajectory angles, 2) testing of the event reconstruction procedure, investigating potential background problems and their impact on the counter performances, and 3) testing the readout electronics. The prototype has been operated with cosmic-ray particles ($CR$) for several months, and tested with $^{12}$C ion beams at various energies at the GSI/Darmstadt facility. \\par This article reports on the results obtained. The counter and its instrumental environment are described in section~\\ref{APPAR}. The readout electronics are presented in section~\\ref{DAQ}, with the data acquisition system used ($DAQ$), the latter for completeness. The analysis procedure is developed in section~\\ref{ANA}, and the results are given and compared with simulation in section~\\ref{RESUL}. The work is summarized and concluded in section~\\ref{DISC}. \\par Some partial results of this work have been reported previously in a few contributions to conferences \\cite{ANTEPROT}. \\par % ",
        "conclusions": "\\label{DISC} The study of a first generation prototype of proximity focussing RICH counter for the AMS experiment reported in this paper has allowed an end-to-end investigation of the technique: Instrumental test of the detector components and electronics,  test of the reconstruction and background rejection algorithm, background measurement, and finally measurement of the counter resolution with different radiator samples using both incident cosmic rays and beam ions with Z$<$6, casting the grounds of the future AMS RICH counter. \\par The above work is being followed by a second generation prototype which incorporates the main features and elements of the final RICH design (flight model). It will be operated using the same instrumental peripheral environment as in the present work. This phase is being undertaken in collaboration between all the institutions involved in the effort on the RICH project \\footnote{INFN Bologna, ISN Grenoble, LIP Lisbon, CIEMAT Madrid, U. Maryland, and IFUNAM Mexico}. \\\\ \\par % \\vspace{0.5cm} {\\bf\\large Acknowledgements.} \\\\ The authors are very indebted to R~.Simon for his invaluable help during the data taking at GSI. They are extremely grateful to M. Yokoyama (Matsushita), J. Favier (LAPP Annecy) and P. Fisher (MIT), for providing aerogel samples, and to B.~Ille (IPN Lyon) for making the set of MWPCs available to the authors. They are also indebted to R. Blanc, T.~Cabanel, G.~Gimon and M.~Marton for their contribution to the detector assembly, to A.~Garrigue, F.~Vezzu, and E.~Perbet, for their contribution to the mechanical study, to J.~Bouvier and O.~Rossetto for their help in the setting up of the electronics, and to Z. Ren for his help on the detector simulation. \\\\ One of the authors (A.M-R.) wishes to acknowledge the ISN hospitality and partial support of CONACYT and DGAPA-UNAM. \\\\ This work was made possible by a dedicated grant from the IN2P3/CNRS.  \\newpage \\begin{center}Appendix: {\\Large \\bf  Refractive index and \\cer radiator thickness} \\end{center} This appendix briefly addresses the issue of the physical variables governing the refractive index and chromatic dispersion of materials. The implication on the thickness of the \\cer radiators is discussed. \\par The relationship between the refractive index  and the phsical of a medium is governed by the Lorentz-Lorenz ({\\it L-L}) law, which can be expressed as \\cite{LLL}: % \\begin{equation} \\frac{n^2-1}{n^2+2}={\\cal N}\\alpha(\\lambda) \\label{LoLoL} \\end{equation} % In this relation, $n$ is the refractive index of the material, {$\\cal N$} the number density of particles in the medium, and $\\alpha(\\lambda)$ the dipole polarizability of the molecules of the medium, i.e., their response function to electromagnetic driving forces. \\par For small values of $(n-1)$ it is straightforward to see that the above can be written as: % \\begin{equation} {n-1}\\approx \\frac{3}{2}{\\cal N}\\alpha(\\lambda) \\label{LoLo2} \\end{equation} % Since ${\\cal N}$ can be expressed in terms of the mass density $\\rho$ of the medium ($\\rho={\\cal\\frac{N}{V} A})$, with ${\\cal A}$ the molar mass of the material, and {$\\cal V$} the Avogadro number), one has the relation of proportionality: \\\\ % \\begin{equation} {n-1}\\approx\\rho\\alpha(\\lambda) \\label{LoLo3} \\end{equation} % This simplified form of the {\\it L-L} equation puts in evidence a few important properties of the refractive index of (transparent) materials: \\\\ 1) - The quantity $(n-1)$ scales with the density $\\rho$ of the material. Therefore, $(n-1)$ will change by approximately 3 orders of magnitude between the gas phase (under atmospheric pressure) and the solid phase for a given element. \\\\ 2) - The dependence of $(n-1)$ on the wave length $\\lambda$ of the incident light is governed by the response function of the molecules of the medium to the corresponding electromagnetic perturbation. The relative variation of $(n-1)$ over a given range of $\\lambda$ is thus given by the relative variation of the molecular response function $\\alpha$: \\\\ % \\begin{equation} \\frac{\\Delta(n-1)}{n-1}\\approx\\frac{\\Delta\\alpha}{\\alpha} \\label{LoLo4} \\end{equation} % \\par Therefore, the scaling law $\\frac{\\Delta(n-1)}{n-1}\\approx constant$ holds rather strictly to within the validity of the approximation for a given material. \\par The derivative of equation \\ref{LoLoL} can be evaluated rigorously however, leading to: \\\\ % \\begin{equation} \\frac{2n}{(n^2+2)(n+1)}\\frac{\\Delta(n-1)}{n-1}=\\frac{\\Delta\\alpha}{\\alpha} \\label{LoLo5} \\end{equation} % The evaluation of the term multiplying the quantity $\\frac{\\Delta(n-1)}{n-1}$ in this relation can be verified to be about constant, close to 0.3 for values of n between 1 and 1.5. The approximation % \\begin{equation} \\frac{\\Delta(n-1)}{n-1}\\approx constant \\end{equation} % is then basically correct, although it is more accurate to use relation \\ref{LoLo5}. \\par\\noindent 3) - It is important to note that the chromatic dispersion of ($n-1$) also scales with the matter density, i.e.: % \\begin{equation} \\Delta(n-1)\\approx\\rho\\Delta\\alpha(\\lambda) \\end{equation} % $\\Delta\\alpha(\\lambda)$ being taken over some relevant range of $\\lambda$. This explains in general why the chromatism of low density materials, like gas or aerogels, is much smaller than that of high density materials like crystals. This explains in particular why it is so for aerogels compared to quartz or fused silica, and it provides a way of estimating the chromatism of the former from the known dispersion law of the latter. \\\\ \\par % {\\bf Thickness of Radiator material} \\\\ The above discussion has straightforward implications for the thickness of the radiator material to be used for a RICH counter. This thickness can be expressed in terms of the \\cer variables. The number of photons radiated is N$_{ph}$=N$_0$Lsin$^2\\theta$, where N$_0$ is the quality factor of the counter \\cite{SIMU}, L the radiator thickness, and $\\theta$ the \\cer angle. One has therefore $L=\\frac{N_{ph}}{N_0 sin^2\\theta}$, or $L\\approx\\frac{N_{ph}}{2N_0<n-1>}$ for small values of $(n-1)$. Using relation \\ref{LoLo3} above: $L\\approx\\frac{N_{ph}}{2N_0\\rho<\\alpha>}$, or % \\begin{equation} \\rho L=\\frac{N_{ph}}{2N_0<\\alpha>} \\label{RHOL} \\end{equation} The quantity $\\rho L$ is the thickness of the radiator in g/cm$^2$. It is seen that this quantity is constant for a given number of photons and for a given material. Although the quality factor can be somewhat different however for different values of $n$, this effect is small for refractive index not too much different, like between 1.02 and 1.1 in silica aerogels. With this restriction, relation \\ref{RHOL} shows that the thickness of material to be used for a given number of photoelectrons does not depend on the mean refractive index of the material with the same molecular structure. For different materials the relation does not hold since the asymptotic value of $\\alpha(\\lambda)$ depends on the value of the first pole of the dispersion law \\cite{LLL}, which can differ by an order of magnitude from one material to another. % %__________________________________________________  \\par"
    },
    "0201/astro-ph0201085_arXiv.txt": {
        "abstract": "This paper  estimates the relative frequency of different types of core-collapse \\se, in terms of the ratio  \\f\\ between the number of type Ib--Ic and  of type II \\se. We estimate \\f\\ independently for all normal and Seyfert galaxies whose   radial velocity  is $\\le$14000 \\kms, and which had at least one \\sa\\ event recorded in the Asiago catalogue from January 1986 to August 2000. We find that the ratio \\f\\ is $\\approx$ 0.23$\\pm$0.05 in normal galaxies. This value is consistent with constant star formation rate and with a Salpeter Initial Mass Function  and average binary rate $\\approx$ 50\\%. On the contrary, Seyfert galaxies exceed the ratio \\f\\ in normal galaxies by a factor $\\approx$ 4  at a confidence level $\\ga 2 \\sigma$. A caveat is that the numbers for Seyferts are still small (6 type Ib-Ic and 6 type II supernovae discovered as yet). Assumed real, this excess of type Ib and Ic with respect to type II supernovae, may indicate a burst of star formation of young age ($\\tau \\simlt$ 20 Myr),  a high incidence of binary systems in the inner regions (r $\\simlt 0.4 $R$_{25}$) of Seyfert galaxies, or a top-loaded mass function. ",
        "introduction": "The relationship between circum-nuclear star formation and non-thermal nuclear activity is as yet poorly understood.  While the possibility that non-thermal activity may be entirely ascribed to massive stars and to supernova events \\cite{terle88,terle94} is challenged by several lines of evidence, it is unlikely that the onset of nuclear activity and strong nuclear or circum-nuclear star formation can be fully unrelated phenomena. Supernov\\ae\\ may not be responsible for the optical spectrum and emission line of Seyfert 1 galaxies, but reprocessed gas of supernova ejections may eventually be accreted by the central black hole, influencing the opacity and hence the radiating properties of the accretion disk which is reputed to be one of the ultimate components of the AGN central engine.  Core-collapse supernov\\ae\\ from massive progenitors, albeit rare events, are diagnostics of recent star formation. The  aim of this paper is to estimate the ratio  between the total number of Ib-Ic \\se\\ type II \\se\\ [\\f=N(Ib/c)/N(II)] discovered in a suitable sample of non-active (hereafter normal) galaxies, as well as to compare it to that of Seyfert galaxies. As we will show through this paper, the ratio \\f\\ reflects metallicity, age, fraction of binary systems, and initial mass function (IMF) shape effects which are probably much different in normal and (at least in some) Seyfert galaxies. ",
        "conclusions": "We found that normal galaxies show \\f$_{norm} \\approx $0.27, a value consistent with the ratio of absolute supernova rates in SNu. With the caveat of our still-limited sample size,  Seyfert galaxies show a peculiar distribution of supernova types, with higher frequency of type Ib/c \\se\\ (\\f$_{Seyf} \\approx~1$) than non-active galaxies. These finding are consistent with a ``normal'' \\sa\\ rate related to secular star formation as far as non-active galaxies are concerned. A scenario where Wolf Rayet stars are produced by the most massive stars ($\\geq$40M$_\\odot$) and by less massive (8M$_\\odot$$\\leq$M$\\leq$40M$_\\odot$) primary stars of binary systems that soon or later undergo mass-transfer, reproduces fairly well the observed ratio, without particular assumptions on the IMF.  A large type Ib/c rate with respect to type II \\se, as found for Seyferts, is a new result. We consider several explanations as possible, namely a lowering of the upper mass limit for type II precursors due to a higher metallicity, a very high {\\em binary fraction}, a top heavy IMF and a young age of the burst of star formation. While metallicity is likely to act in the observed direction, because it enhances the mass-loss rate, the amplitude of the observed effect  may not be explained.  Even if the \\f\\ value found for Seyfert galaxies  needs confirmation with better statistics, it is interesting to note that an high \\f\\ could be associated with the enhanced SFR in the circum-nuclear and regions of galaxies hosting Seyfert nuclei, especially of type 2 nuclei \\cite{dd95}. A young age of the burst of star formation -- probably triggered by interaction with a massive companion galaxy \\cite{dultzin99} -- is a likely cause of the high \\f\\ observed among Seyfert galaxies, as it requires the minor number of assumptions. This interpretation is also consistent with an evolutionary sequence leading from IR-luminous Starbursts, to Seyfert 2 and eventually to Seyfert 1 galaxies, with the youngest Seyfert 1 being the ones of lowest intrinsic power and/or the ones most heavily obscured \\cite{gu01}.   \\bigskip AB thanks the {\\sl Osservatorio di Collurania} at Teramo (Italy)  for their hospitality. Two of us  acknowledge financial support by the Italian Ministry for University and Research (MURST) under grant Cofin 00-02-007 (PM) and Cofin 9902763542\\_004 (AB)."
    },
    "0201/astro-ph0201521.txt": {
        "abstract": "We review recent results on the nonlinear development of thermal instability (TI) in the context of the turbulent atomic interstellar medium (ISM), in which correlated density and velocity fluctuations are present, as well as forces other than the thermal pressure gradient. First, we present a brief summary of the linear theory, remarking that, in the atomic ISM, the condensation mode is unstable but the wave mode is stable at small scales. Next, we revisit the growth of isolated entropy perturbations in initially unstable gas, as a function of the ratio of the cooling to the dynamical crossing times $\\eta$. %, by %means of high-resolution one-dimensional numerical simulations. %We quantify the magnitude %of the velocities developing The time for the dynamical transient state to subside ranges from 4 to 30 Myr for initial density perturbations of $20\\%$ and sizes 3 to 75 pc. When $\\eta \\ll 1$, the condensation produces locally supersonic motions and a shock propagates off the condensation, bringing the surrounding medium out of thermal equilibrium. %By the time the %condensations have formed, a substantial fraction of the mass is still %traversing the unstable range. Third, we consider the evolution of {\\it velocity} perturbations, maintained by a random forcing, representing turbulent energy injection to the ISM from stellar sources. These perturbations correspond to the wave mode, and are stable at moderate amplitudes and small scales, as confirmed numerically. %Numerical simulations show that %these perturbations do not grow over %arbitrarily long times if the forcing is applied at small enough scales %that $\\eta \\gg 1$, and causes a moderate rms Mach number ($\\gtrsim 0.3$). %at those scales. %From the two sets of results, we suggest %that, when only TI and random turbulent motions are considered, no %sharp phase transitions between the cold and warm phases of TI should be %expected, and that a fraction of the gas should be in the intermediate, %``unstable'' range.  We then consider the behavior of magnetic pressure in turbulent regimes. Various observational and numerical results suggest that the magnetic pressure does not correlate well with density at low and intermediate densities. We propose that this is a consequence of the slow and fast modes of nonlinear MHD waves being characterized by different scalings of the magnetic field strength versus density. This lack of correlation suggests that, in fully turbulent regimes, the magnetic field may not be a very efficient source of pressure, and that polytropic descriptions of magnetic pressure are probably not adequate.  Finally, we discuss simulations of the ISM (and resolution issues) tailored to investigate the possible existence of significant amounts of gas in the ``lukewarm'' temperature range between the warm and cold stable phases. The mass fraction in this range increases, and the phase segregation decreases, as smaller scales are considered. We attribute this to two facts: the enhanced stability of moderate, adiabatic-like velocity fluctuations with $\\eta \\gg 1$ and the recycling of gas from the dense to the diffuse phase by stellar energy injection. %This appears to be a consequence of the moderate turbulence %found in the intercloud medium, which can induce much more stable and %out-of-thermal equilibrium adiabatic-like perturbations. Moreover, the magnetic field is not strongly turbulent there, possibly providing  additional stability. We conclude by suggesting that the gas with unstable temperatures can be observationally distinguished through simultaneous determination of two of its thermodynamic variables. ",
        "introduction": "\\label{sec:intro}  The fact that the neutral atomic interstellar medium (ISM) is most likely thermally bistable \\cite{Pik68,FGH69,Wol95} has had a great impact on our picture of interstellar structure formation. Indeed, in two of the most influential models of the ISM to date, the famous two- and three-phase models of the ISM of Field, Goldsmith \\& Habing \\cite{FGH69} and McKee \\& Ostriker \\cite{MO}, the concepts of thermal and pressure equilibria played a fundamental role, so that distinct {\\it phases} (thermodynamic regimes with different density and temperature, but the same pressure) were predicted to coexist in pressure equilibrium. These phases correspond to stable thermal-equilibrium (i.e., heating-cooling balance) temperature regimes, and are separated by unstable regimes that, in those models, were therefore not expected to be present in the ISM. An opposite view was taken in the so-called time-dependent model of the ISM of Gerola, Kafatos \\& McCray \\cite{GKM74}, which was presented as an alternative to the pressure equilibrium two-phase model, and which made radically different assumptions: a constant density in the presence of stochastic, local heating events that caused strong local fluctuations of pressure and temperature, because the cooling and recombination times are comparable or shorter than the time between successive exposures of a given gas parcel to those heating events. This model predicted that significant amounts of gas should be in the unstable range, as they cooled after the transient heating events. More recently, Lioure \\& Chi\\'eze \\cite{LC90} have considered models with a continuous recycling of gas among the various gas phases due to stellar energy injection, also concluding that significant amounts of gas should populate the unstable temperature range in the ISM. Note that the three-phase model \\cite{MO} did consider the existence of local fluctuations in the pressure, although it was still based on the premise of ``rough pressure balance''.  Nevertheless, both the equilibrium and the time-dependent models omitted a number of important aspects in the ISM budget. The multiphase equilibrium models essentially neglected the possibility of large pressure fluctuations in the ISM. The time dependent model instead included this possibility as a fundamental premise, but neglected the fact that such pressure fluctuations should induce motions, which should in general be turbulent (i.e., spanning a wide range of scales), and in turn cause strong density fluctuations \\cite{Elm93,BVS99}. Moreover, both the time-dependent and the three-phase models omitted other important agents of the ISM, such as magnetic fields, rotation, and cosmic rays. Elmegreen \\cite{Elm91,Elm94} performed a combined instability analysis including self-gravity, cooling and heating, and magnetic fields, but the effects of turbulence, which is an inherently nonlinear phenomenon, can only be dealt with by means of numerical simulations of the gas dynamics in the Galactic disk in the presence of thermal instability (TI). The role of the turbulent motions may be crucial. In fact, realistic cloud/intercloud structure has been reported in models incorporating turbulence from stellar-like driving and cooling, but not necessarily a thermally bistable regime \\cite{BL80,CP85,CB88,RB95,VPP95,PVP95,GP99}. Similar results have been reported for pressureless (Burgers-like) models with stellar driving \\cite{SC99,CS01}, and for simulations of interacting nonlinear MHD waves \\cite{Elm97}. Thus, it is important to investigate the role of TI in determining the distribution of the physical variables (density, temperature, velocity) of the flow, and, in particular, the degree to which phase segregation, as was proposed in the multiphase models, is realized, in the context of a turbulent ISM with multiple sources of turbulent energy at a variety of scales, such as stellar winds, supernova explosions, spiral arm passage, magnetorotational instabilities \\cite{SB99}, etc., besides TI.  Although the nonlinear development of TI has been studied extensively for decades now (e.g., \\cite{Gol70,SMS72,Muf74,Muf75,OMB82,Sas88,DBS88,ML89,ML90,BMM90,Kri90,KSFR90,KLR00,BL00}), only recent work has started to investigate the interplay between TI and the turbulent nature of the ISM, such as, for example, the triggering of TI by external compressions (\\cite{HP99,HP00,KI00}) and the possibility that the TI itself may contribute to the generation of turbulence in the ISM \\cite{WSK00,KI01,KN02}. %The latter effect occurs because %long-wavelength perturbations can develop large (supersonic) %velocities and strong shocks as they condense under TI %\\cite{Fie65,MS87,Bal95,Mee96}. In flows in more than one dimension, %dynamical instabilities may thus generate turbulence, as those works %suggest.  However, the fact that additional energy sources feeding ISM turbulence besides TI itself, such as stellar energy injection, or large-scale gravitational or magnetic instabilities, has additional implications. First, the very presence of strong motions implies that transport (advection) should be important, while traditionally conductive processes have received more attention (e.g., \\cite{ZP69,BM90,Mee96}). Second, these transport processes may imply the existence of constant fractions of gas transiting through the unstable regime, and erase, to some extent, the phase segregation expected in multiphase models. These expectations are furthered by several observational studies (e.g., \\cite{DST77,KSG85,SF95,FS97,Hei01}) that have suggested that the fraction of gas in the unstable range between the cold and warm phases of the atomic ISM is substantial.  Another important issue to consider is the fact that the ISM is magnetized, which suggests the possibility that turbulent magnetic pressure may supplement thermal pressure and somehow counteract TI. The condensation process in a magnetized medium has been studied by several workers (see, e.g. \\cite{Fie65,Elm97,OMB82,Loe90,HP00}; see also the references given in \\cite{HP00}), concluding that, although condensation can be inhibited under some circumstances, it is in general possible. However, those studies have not considered the case of TI developing in an externally-driven turbulent medium, except for \\cite{HP00}.  In this paper we review recent work and present new results concerning the interplay between TI and turbulence in the warm and cool ISM. First, we review the main aspects of the instability in \\S \\ref{sec:quali}. Then, in \\S \\ref{sec:pure_TI}, we revisit the growth of isolated density fluctuations, focusing in particular on the late stages and the state of the gas surrounding the condensation. %quantify, by means of %high-resolution numerical simulations in one dimension (1D), the %magnitudes of the motions produced by the nonlinear development of TI in %the atomic medium, and their associated durations, as %a function of the ratio of the cooling to sound crossing times, %hereafter denoted $\\eta$. Next we discuss the development or suppression of growth in the presence of random velocity fluctuations, stressing that these probably constitute the most common way of inducing density fluctuations in the ISM. We then briefly discuss the nature of magnetic pressure on turbulent media (\\S \\ref{sec:magn_pres}) and its dependence on density. In \\S \\ref{sec:full_ISM}, we discuss the role of TI in numerical models of the ISM in the presence of magnetic fields, the Coriolis force, modeled star formation, self-gravity and TI, aiming at determining the fraction of unstable gas, and at interpreting the results in the light of the previous sections. We include an extensive discussion of numerical tests to maximize the reliability of the results. Finally, we summarize the results and mention a number of implications in \\S \\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  In this paper we have reviewed results from various studies aiming at understanding the role of TI in the turbulent atomic ISM, and the behavior of the magnetic pressure in the fully turbulent case. The motivation has been twofold. On the one hand, the classic multi-phase models of the ISM have neglected the implications of the ISM being turbulent, and it is thus important to assess the consequences of advection on the thermal and spatial structure of this medium. On the other hand, observations have often suggested the presence of gas with temperatures in the thermally unstable range, in apparent contradiction with the multi-phase models.  We first reviewed the classic instability analysis of Field \\cite{Fie65}, emphasizing the different behavior of long- and short-wavelength perturbations (for which the ratio $\\eta$ of the cooling [$\\tc$] to the sound crossing [$\\ts$] time is respectively small and large), and of entropy and isentropic perturbations (which trigger the condensation and the wave modes, respectively). We pointed out that, while much study has been devoted to isobaric entropy perturbations, real-world fluctuations in the ISM are produced through velocity fluctuations which, in the small-scale limit, belong to the isentropic kind, and are therefore stable to first order at small scales. We also briefly reviewed the magnetic case, in which the presence of a uniform magnetic field can stabilize perturbations with wavenumbers perpendicular to it.  We then reviewed results on the nonlinear stages of evolution of isobaric entropy perturbations, focusing on those that have quantified the magnitude of the speeds developed and the times required for completing the condensation process as a function of the parameter $\\eta$. For presently-accepted values of the heating and cooling rates \\cite{Wol95}, large-scale initial perturbations ($\\gtrsim 15$ pc, $\\eta \\lesssim 0.2$) develop supersonic speeds, require times $\\gtrsim 10$ Myr to complete the condensation process, and end up with densities and pressures above the thermal equilibrium value due to the ram pressure of the still infalling gas. Those times are long compared with typical times between successive external shock passages, and star formation time scales. We thus concluded that clouds formed from perturbations of such sizes (although the resulting cloud has a size $\\sim 1$ pc) are unlikely to exist in thermal pressure equilibrium with their surroundings. Initial perturbations of sizes $\\lesssim 3$ pc, on the other hand, require times $\\sim 4$ Myr to complete their evolution and do not generate supersonic speeds, thus reaching a more quiescent final state, and adhering better to the paradigm of thermal-pressure bounded clouds at the end of their evolution, although, by the time the cloud has formed, accretion is still occurring, so that the clouds are bounded by weak accretion fronts, rather than contact discontinuities. Furthermore, the gas still accreting is necessarily in the unstable temperature range, although it is not in thermal equilibrium; instead, it has a ``regular'' pressure behavior ($P$ increases with $\\rho$), and thus it is not prone to further fragmentation.  We then described the evolution of perturbations induced by turbulent random forcing. In this case the crossing time entering $\\eta$ should be taken as the minimum of the sound and the turbulent ($\\tu$) crossing times. Thus, small-scale ($\\sim 0.3$ pc) {\\it velocity} fluctuations are quasi-isobaric at very small amplitudes, because in this case $\\tu > \\tc > \\ts$ so that the flow can cool in response to the velocity perturbation. As the perturbation amplitude is increased, so that $\\tc > \\ts > \\tu$, the situation changes because now the density is driven by the turbulent velocity rather than by sound waves, and the perturbations become quasi-adiabatic in character, {\\it becoming stable}. We empirically found this to occur roughly when the rms Mach number $\\gtrsim 0.3$. Finally, however, if the perturbation amplitude becomes very large, then the density increment induced by it becomes nonlinear and accelerates the cooling rate, effectively causing $\\eta <1$. In this case, velocity fluctuations trigger the condensation mode, which is again unstable, and cause condensations.  Thus, we reached the important conclusion that small-scale fluctuations behave very differently when they are entropy perturbations (caused, for example, by local variations in the heating or cooling rates) and when they are adiabatic (caused by velocity fluctuations), being unstable (and with the fastest growth rates) in the former case, but linearly stable in the latter.  We then considered the magnetic field as an additional source of pressure in the ISM, confirming earlier results that at low and intermediate densities the magnetic pressure is strongly decorrelated from density in fully turbulent cases (large field fluctuations), and proposed an interpretation of this phenomenon in terms of the scaling of $B^2$ with density for the slow and fast modes of simple nonlinear MHD waves. The decorrelation between magnetic pressure and density has several implications, among which is that the magnetic field probably is ineffective in supplementing thermal pressure in highly turbulent, thermally unstable conditions, and that it is probably inadequate to model magnetic pressure by means of an equivalent polytropic behavior in the fully turbulent case.  Finally, we discussed results from simulations of the ISM in more than one dimension at large and intermediate scales and at various resolutions. To this end, we first performed a detailed study of the competition between numerical diffusivities and the growth of TI, finding that even when the diffusivities (especially the mass diffusion, which is necessary numerically) are confined to the smallest scales on the numerical grid, they can push the smallest unstable scale (the ``Field'' length $\\lf$) to relatively large scales in the simulations, especially for small physical simulation sizes, because $\\lf$ scales as $\\lambda_K^{1/2}$, where $\\lambda_K$ is the diffusive scale.  With this information, we discussed the fact that many ISM simulations suggest that the basic structure does not depend sensitively on whether TI is present, as long as there are turbulent motions driven by stellar-like sources (that imply recycling of gas from the cold to the warm phase), and that significant fractions of the gas mass (15-50\\%) appear to be in the unstable regime. This appears to be a consequence of the fact that the diffuse medium is in a moderately turbulent state, so that a) the fluctuations there have a regular pressure gradient and b) the magnetic field is not strongly turbulent, and therefore may cause additional stability. Of course, when the relatively quiescent intercloud medium is hit by a strong shock from, say, a supernova remnant, then TI can be rapidly induced, as in the studies by Hennebelle \\& P\\'erault \\cite{HP99,HP00} and Koyama \\& Inutsuka \\cite{KI01}.  A final remark of interest is that it may be possible to determine observationally whether the gas seen at unstable temperatures corresponds to the out-of-thermal-equilibrium gas observed in the simulations by either a) simultaneously determing two of its thermodynamic variables, or b) comparing directly observed cooling rates (e.g., fine structure lines) with theoretical estimates of the heating rate (e.g., photoelectric heating) in specific regions (C.\\ Heiles, private communication). If this is confirmed, then it would provide strong evidence that turbulent motions populate all regions of the thermodynamic variable space, preventing a sharp segregation of the atomic ISM into the stable phases of TI.  \\bigskip We have greatly benefitted from exchanges with C. Heiles, P. Hennebelle, H. Koyama, J. Scalo and E. Zweibel. The report from an anonymous referee prompted much improvement of the paper and led us to the study of numerical damping of the growth rates. This work has received partial financial support from CONACYT grant 27752-E, from the French national program PCMI, and from the conference organizers to E.V.-S. We have made extensive use of NASA's Astrophysics Data System Abstract Service.    %"
    },
    "0201/astro-ph0201200_arXiv.txt": {
        "abstract": "{We discuss new photometric data collected on the $\\gamma$ Dor variables HD~224945 and HD~224638. Multiperiodicity was detected in both stars, thanks to the clear spectral window of a multisite campaign that involved five observatories. HD~224945 shows the shortest period among the $\\gamma$ Dor stars, i.e., 0.3330~d. The pulsation behaviour is very different: HD~224945 displays a set of frequencies spread over an interval much wider than that of HD~224638. We clearly found evidence for amplitude variations in the excited modes by comparing data from different years. HD~224945 and HD~224638 are among the best examples of $\\gamma$ Dor stars that show multimode pulsations, which make them very interesting from an asteroseismological point of view. ",
        "introduction": "The variability of HD 224638$\\equiv$BT Psc ($V$=7.5, F1~V) and HD 224945$\\equiv$BU Psc ($V$=6.93, F0~V) was announced by Mantegazza \\& Poretti (1991), as a by--product  of the monitoring of the $\\delta$ Sct star HD 224639$\\equiv$BH Psc. Both stars had been used as comparison stars in the first observing run devoted to BH Psc. They increased the number of known $F$--type stars located close to the low--temperature edge of the Cepheid instability strip which exhibit small amplitude variability on time scales of several hours, usually longer than the length of a night of observation and therefore easily detectable only when used as comparison stars for short--period variables. At that time, the debate on the nature of these variations was divided between spot activity (observed periods as rotational periods) and pulsation (nonradial $g$--modes). Mantegazza et al. (1994, hereinafter Paper~I) tried to explain the complicated light behaviour of HD 224639 and HD 224945 in the simplest way possible, by means of periodicities shaped as double-- or triple--wave curves.  Balona et al. (1994) reported on the multiperiodicity of $\\gamma$ Dor, giving decisive evidence in favour of pulsation. Also, as a result of combined photometry and line profile variations for 9 Aurigae (Krisciunas et al. 1995) and $\\gamma$ Dor (Balona et al. 1996), plus the preliminary results on HD 224945 (Poretti et al. 1996), the  hypothesis of variability caused by $g$--modes  achieved a wide consensus. The properties of this new  class of variable stars were delineated step--by--step through observational efforts and have been summarized by Kaye et al. (1999), while the problem of the driving mechanism and the excitation of $g$--modes constitutes the target of continuing theoretical investigations. As a result, $\\gamma$ Dor variables are now  considered intermediate between $A-F$ pulsators and solar--type stars, finding  a special place  in the main programmes of asteroseismologic missions such as {\\sc corot} and {\\sc mons}.  On the basis of this progress, it is worth investigating in more detail the pulsational behaviour of HD 224639 and HD 224945. Here we present the results of a multisite campaign carried out in October 1995. ",
        "conclusions": "The solution of the light curves of HD~224638 and HD~224945 has been deduced only on the basis of a multisite campaign, since the true peaks could not be recognized if a $\\pm1~$\\cds effect is present in the spectral window. We did our best to identify the correct excited modes, but the more important characteristic of these two stars is the strong multiperiodicity itself, independent of the exact frequency values. In our opinion,  HD~224638 and HD~224945 provide two of the best examples of features typical of  multiperiodicity among $\\gamma$ Dor variables: different sets of frequency content, amplitude variations, disappearing terms, close doublets of frequencies.  Also, after detecting five or six terms, the rms scatter is larger than the observational error. Therefore, residual signal is hidden in the noise. Such extreme multiperiodicity can be clearly ascribed to pulsation, as stellar activity is not able to generate it.   The differences in the frequency and amplitude ranges of HD~224638 and HD~224945 (clearly visible in Fig.~\\ref{lc}) remind us of the unpredictable frequency content of $\\delta$ Sct stars, where the selection mechanism among all the possible modes seems to be different from one star to the next (Poretti 2000). The amplitude variability of the excited modes is another point of similarity between $\\gamma$ Dor and $\\delta$ Sct stars. Other than multiperiodicity, HD~224945 and HD~224638 provide the best examples of amplitude variability among $\\gamma$ Dor stars, following the three campaigns carried out on these stars. This observational evidence suggests strategies for  asteroseismological space missions: a long observing run (or two separate runs) may be very helpful in detecting more terms (excited at different levels at different times) or in studying damping effects."
    },
    "0201/astro-ph0201493_arXiv.txt": {
        "abstract": "Efforts to understand the deviation of the L--T relation from a simple scaling law valid for clusters and groups have triggered a number of interesting studies on the subject. Techniques and approaches differ widely but most works agree on the important role played by gas cooling and heating sources like AGNs and SNe. Observations set useful constraints on the evolution of the intracluster medium (ICM): a 100KeV/cm$^2$ entropy floor in the core of groups and about 5--15\\% of baryons being converted into stars. However, essential details like the nature of the dominant heating mechanism and the quantitative importance of cooling still need to be addressed. I suggest that a new generation of high resolution N-body simulations and a quantitative comparison of results between different approaches is required to improve results and increase our understanding of the problem. ",
        "introduction": "The first attempt to model the ICM in the framework of the hierarchical scenario assumed its thermodynamical properties to be entirely determined by gravitational processes, like adiabatic compression during collapse and shock heating (Kaiser 1986). If no characteristic scales are present in the underlying cosmology (i.e., Einstein--de-Sitter cosmology and power--law shape for the power spectrum of density perturbations), this model should predict hot gas within rich clusters to look the same as within poor groups, since gravity in itself does not have characteristic scales. Under the assumptions of emissivity dominated by free--free bremsstrahlung and of hydrostatic equilibrium of the gas, this model predicts $L_X\\propto T^2(1+z)^{3/2}$ for the shape and evolution of the relation between $X$--ray luminosity and ICM temperature (also Eke, Navarro \\& Frenk 1998). Furthermore, if we define the gas entropy as $S=T/n_e^{2/3}$ ($n_e$: electron number density; e.g., Eke et al. 1998), then the self--similar ICM has $S\\propto T(1+z)$. ",
        "conclusions": ""
    },
    "0201/astro-ph0201170_arXiv.txt": {
        "abstract": "{ The general PDE governing linear, adiabatic, nonraradial oscillations in a spherical, differentially and slowly rotating non-magnetic star is derived. This equation describes mainly low-frequency and high-degree $g$-modes, convective $g$-modes, and rotational Rossby-like vorticity modes and their mutual interaction for arbitrarily given radial and latitudinal gradients of the rotation rate. Applying to this equation the `traditional approximation' of geophysics results in a separation into radial- and angular-dependent parts of the physical variables, each of which is described by an ODE. The condition for the applicability of the traditional approximation is discussed. The angular parts of the eigenfunctions are described by Laplace's tidal equation generalized here to take into account differential rotation. From a qualitative analysis of Laplace's tidal equation the sufficient condition for the formation of the dynamic shear latitudinal Kelvin-Helmholtz instability (LKHI) is obtained. A small rotation gradient causes LKHI of prograde waves (seen in the rotating frame), while strong gradients are responsible for retrograde LKHI. The value of the latitudinal rotation gradient has a lower limit, below which LKHI disappears. The LKHI result is applied to real solar helioseismology rotation data. It is shown that the $m=1$ mode ($m$ = azimuthal wave number) instability can develop. This global instability takes place in the whole envelope of the Sun, including the greatest part of the tachocline, in radial direction and at almost all latitudes in horizontal direction. The exact solutions of Laplace's equation for low frequencies and rigid rotation are obtained. There exists only a retrograde wave spectrum in this ideal case. The modes are subdivided into two branches: fast and slow modes. The long fast waves carry energy opposite to the rotation direction, while the shorter slow-mode group velocity is in the azimuthal plane along the direction of rotation. The eigenfuncions are expressed by Jacobi's polynomials which are polynomials of higher order than the Legendre's for spherical harmonics. The solar 22-year mode spectrum is calculated. It is shown that the slow 22-year modes are concentrated around the equator, while the fast modes are around the poles. The band of latitude where the mode energy is concentrated is narrow, and the spatial place of these band depends on the wave numbers ($l, m$). % ",
        "introduction": "In a recent paper Dzhalilov et al. (2001; paper~1) investigated which lowest-frequency eigenoscillations can occur in the real Sun, moreover, which role they play in redistributing angular momentum and causing solar activity. We found that such waves could only be differential rotation Rossby-like vorticity modes. However, the general nonradial pulsation theory adopted from stellar rotation has some difficulities. For slow rotation, when the sphericity of the star is violated not seriously, the degeneracy of the high-frequency spherical $p$-  and $g$-modes with respect to the azimuthal number $m$ is abandoned by rotation (Unno et al. 1989). Independent of the spherical modes non-rotating toroidal flows (called `trivial' modes with a zero frequency) become quasi-toroidal with rotation (called $r$-modes with a nonzero frequency; Ledoux 1951; Papaloizou \\& Pringle 1978; Provost et al. 1981;  Smeyers et al. 1981; Wolff 1998). Although rotation abandons the degeneracy of the modes, it also couples the modes with the same azimuthal order, and this makes the problem more difficult. For the high-frequency modes ($\\en_{\\sR}=\\omega/2\\Omega\\ge 1$, where  $\\omega$ and $\\Omega$ are the angular frequencies of oscillations and of stellar rotation, respectively) this difficulty is resolved more or less successfully. For this case the small perturbation rotation theory is applied, in  which the eigenfunctions are represented by power series, the angular parts of which are expressed by spherical harmonic functions $Y^m_l$ (Unno et al. 1989). These power series are well truncated, unless $\\en_{\\sR}<1$, when the role of Coriolis  force is increasing.  Namely the low-frequency instabilities are discovered in most pulsating stars (Cox 1980; Unno et al. 1989). Rotation couples strongly together the high-order $g$, the convective $g$, and the $r$-modes with $\\en_{\\sR}<1$ and with the same $m$, but different $l$ (Lee \\& Saio 1986). Generally the matrix of the coupling coefficients to be determined is singular (e.g. Townsend 1997). In all papers on the eigenvalue problem of nonradially pulsating stars, there exists a `truncation problem' for the serial eigenfunctions, the angular parts of which are represented by spherical harmonics (e.g. Lee \\& Saio 1997; Clement 1998).  The governing partial differential equations (PDEs) of the eigenoscillations of rotating stars are complicated from the point of view of the mathematical treatment, even if the motions  are adiabatic. This difficulty arises because in spherical geometry an eigenvalue problem with a singular boundary condition has to be solved. These equations are simplified considerably to neglect the tangential components of the angular velocity $\\vec\\Omega$ in the low-frequency case $\\en_{\\sR}<1$ (this means that the motion caused by the Coriolis force is primarily horizontal). This limitation widely used in geophysical hydrodynamics (e.g. Eckart 1960) is called `traditional approximation' and has been used first by Laplace (1778) to study tidal waves (Lindzen \\& Chapman 1969). Laplace's equation (or the traditional approximation) for $\\en_{\\sR}<1$ is applicable to the stellar case too. The main advantage of this approximation is that it decomposes the initial system of equations into a pair of ordinary differential equations (ODEs) (e.g. Lindzen \\& Chapman 1969; Berthomieu et al. 1978; Bildsten et al. 1996; Lee \\& Saio 1997). The angular parts of the eigenfunctions are described by Laplace's tidal equation. Solving this equation numerically by using a relaxation method, Lee and Saio (1997) first avoided the representation of the solutions by $Y_l^m(cos\\theta)$ functions for the $\\Omega=$const case, and they had no problem with the truncation of the series.  In the present work for the non-magnetic and non-convective cases we receive one PDE in spherical geometry for the adiabatic pressure oscillations in the differentially rotating star ($\\Omega=\\Omega(r, \\theta)$) with arbitrary spatial gradients of rotation (Sect.~2). This general equation is split into the $\\theta$- and $r$-component ODEs, if the traditional approximation is applied (Sect.~3). The $\\theta$-component equation is Laplace's tidal equation generalized for the differentially rotating case. In Sect.~4 we analyse more qualitatively this equation. We find the general condition for the shear instability due to differential rotation in latitude. We find that the smallest rotation gradient is responsible for the prograde (seen in the rotating frame) vorticity wave instability, while a stronger gradient causes the retrograde wave instability. For solar data (small rotation gradients) the $m=1$ prograde mode instability is possible (Sect.~4.4). The possible existence of such a global horizontal shear instability on the Sun has been investigated by Watson (1981) and Gilman \\& Fox (1997), that of shear and other dynamic instabilities and of thermal-type instabilities in stars as well by Knobloch \\& Spruit (1982) and others. Laplace's tidal equation for low frequencies in the rigid-rotation case is investigated in detail in Sect.~5. It is shown that the eigenfunctions are defined by Jacobi's polynomials which are of higher order than the Legendre's. ",
        "conclusions": "In the present work we have derived the general PDE governing non-radial, adiabatic, long-period (with respect to the rotation period), linear oscillations of a slowly and differentially rotating star. This general equation includes all the high-order $g$-modes and all possible hybrids of rotation modes as well as their mutual interaction. The geophysical `traditional approximation' considerably simplifies this general equation, and we get two ODEs for the $r$- and $\\theta$-components instead of one with arbitrary gradients of rotation $\\Omega(r,\\theta)$. We have received a more stringent condition for the applicability of this approximation to the pulsation of stars. Only for very low frequencies this restriction is the same as that of the standard case.  The $\\theta$-equation is Laplace's equation generalized to the latitudinal differential rotation. Without solving this equation qualitatively we found the exact condition for the appearance of a global instability. This instability is driven by the latitudinal shear, it is not influenced by buoyancy. We call that a `latitudinal Kelvin-Helmholtz instability' (LKHI). The appearance of LKHI strongly depends on the Rossby number (the ratio of rotation period and period of motion), on the azimuthal wave numbers and on the latitudinal rotation gradients. Very large gradients produce retrograde waves (seen in the rotating frame), while a slower rotation gradient is responsible for prograde mode LKHI. The rotation gradient has a lower boundary below which LKHI is not possible for any Rossby number or azimuthal number $m$.  We have applied the LKHI condition to the helioseismological data of the Sun. Here a global LKHI is possible for the $m=1$ mode at practically all latitudes. Radially the LKHI is extended from the greatest part of the tachocline up to the photosphere. The LKHI for the Sun was first obtained by Watson (1981). According to his results the instability is possible only at photosphere layers. Later Gilman \\& Fox (1997) have shown that such an instability is possible in the tachocline too, if strong toroidal magnetic fields are included. Our results show that the instability of the $m=1$ modes and other modes is possible without magnetic fields, in contradiction to Gilman \\& Fox (1997). This difference is probably connected with the incompleteness of the equations used by Watson (1981) and by Gilman \\& Fox (1997); their equations are two-dimensional only.  The exact solutions of Laplace's tidal equation for lower frequencies are expressed by Jacobi's polynomials. Just for lower frequencies the numerical calculations of stellar pulsation analyses meet great problems, when looking for the eigenfunction as infinite series of Legendre functions. The eigenfunctions, defined by higher-order polynomials of Jacobi, cannot be expressed by convergent series of associated Legendre functions. Every Legendre function is a particular case of a Jacobi polynomial.  It has been shown here that the retrograde (slow and fast) modes with high surface wave numbers ($l,m$) are energetically concentrated in narrow bands of latitudes. This analysis was done for the 22-year modes as an example. Such a concentration of mode energy in a narrow spatial area makes such modes vulnerable to different instability mechanisms such as the $\\en$--mechanism considered in Paper~1.  \\appendix  \\begin{table*} \\centering \\caption{Results for the 22-year period: frequency deviations $\\delta=\\omega_{22}-\\omega$ for permitted quantum numbers $(l, |m|)$} \\begin{tabular}[]{rrrrr|@{\\ \\ }|rrrrr|@{\\ \\ }|rrrrr} \\hline $l$  & $m$ \\ & $\\delta\\ \\ \\ $ &$m$ \\ & $\\delta\\ $ & $l$  & $m$ \\ & $\\delta\\ \\ \\ $ &$m$ \\ & $\\delta\\ $ &$l$ & $m$ \\ & $\\delta\\ \\ \\ $ & $m$ \\ & $\\delta\\ \\ \\ $ \\\\ & fast& (nHz) &  slow & (nHz)  &   & fast& (nHz) &  slow & (nHz)  & &  fast & (nHz)  & slow & (nHz) \\\\ \\hline\\hline 11 &   1 &  0.055 & 575 & 0.000 & 33 &  10 & -0.021 & 478 & 0.000 & 55 &  35 &-0.008 &   365 &-0.001\\\\ 12 &   1 & 0.250  & 571 & 0.001 & 34 &  10 & 0.051  & 474 & 0.001 & 56 &  36 & 0.010  & 360 & 0.001 \\\\ 13 &   1 & 0.406 & 567 & 0.001 & 35 &  11 & 0.020 & 469 & 0.000 & 57 &  38 & 0.007  & 354 & 0.001 \\\\ 14 &   2 &-0.265 & 562 &-0.001 & 36 &  12 &-0.003 & 464 & 0.000 & 58 &  40 & 0.007 & 348 & 0.001 \\\\ 15 &   2 &-0.073 & 558 & 0.000 & 37 &  13 &-0.021 & 459 &-0.001 & 59 &  43 &-0.008 & 341 &-0.001 \\\\ 16 &   2 & 0.089 & 554 & 0.000 & 38 &  14 &-0.033 & 454 &-0.001 & 60 &  45 &-0.004 & 335 & 0.000 \\\\ 17 &   2 & 0.226 & 550 & 0.001 & 39 &  14 & 0.029 & 450 & 0.001 & 61 &  47 & 0.002 & 329 & 0.000 \\\\ 18 &   3 &-0.128 & 545 &-0.001 & 40 &  15 & 0.020 & 445 & 0.001 & 62 &  50 &-0.005 & 322 &-0.001 \\\\ 19 &   3 & 0.013 & 541 & 0.000 & 41 &  16 & 0.014 & 440 & 0.001 & 63 &  52 & 0.004 & 316 & 0.001 \\\\ 20 &   3 & 0.136 & 537 & 0.001 & 42 &  17 & 0.012 & 435 & 0.000 & 64 &  55 & 0.001 & 309 & 0.000 \\\\ 21 &   4 &-0.091 & 532 &-0.001 & 43 &  18 & 0.012 & 430 & 0.000 & 65 &  58 & 0.002 & 302 & 0.000 \\\\ 22 &   4 & 0.028 & 528 & 0.000 & 44 &  19 & 0.014 & 425 & 0.001 & 66 &  61 & 0.004 & 295 & 0.001 \\\\ 23 &   4 & 0.134 & 524 & 0.001 & 45 &  20 & 0.018 & 420 & 0.001 & 67 &  65 &-0.001 & 287 & 0.000 \\\\ 24 &   5 &-0.021 & 519 & 0.000 & 46 &  22 &-0.019 & 414 &-0.001 & 68 &  69 &-0.002 & 279 &-0.001 \\\\ 25 &   5 & 0.079 & 515 & 0.001 & 47 &  23 &-0.010 & 409 &-0.001 & 69 &  73 &-0.001 & 271 & 0.000 \\\\ 26 &   6 &-0.035 & 510 & 0.000 & 48 &  24 & 0.001 & 404 & 0.000 & 70 &  77 & 0.002 & 263 & 0.000 \\\\ 27 &   6 & 0.058 & 506 & 0.001 & 49 &  25 & 0.011 & 399 & 0.001 & 71 &  82 & 0.001 & 254 & 0.000 \\\\ 28 &   7 &-0.029 & 501 & 0.000 & 50 &  27 &-0.009 & 393 &-0.001 & 72 &  88 & 0.000 & 244 & 0.000 \\\\ 29 &  7  & 0.057 & 497 & 0.001 & 51 &  28 & 0.005 & 388 & 0.000 & 73 &  94 & 0.001 & 234 & 0.001 \\\\ 30 &   8 &-0.009 & 492 & 0.000 & 52 &  30 &-0.009 & 382 &-0.001 & 74 & 102 & 0.000 & 222 & 0.000 \\\\ 31 &   9 &-0.060 & 487 &-0.001 & 53 &  31 & 0.006 & 377 & 0.001 & 75 & 112 &-0.001 & 208 & 0.000 \\\\ 32 &   9 & 0.019 & 483 & 0.000 & 54 &  33 &-0.002 & 371 & 0.000 & 76 & 125 & 0.000 & 191 & 0.000 \\\\ \\hline \\end{tabular} \\end{table*}"
    },
    "0201/astro-ph0201346_arXiv.txt": {
        "abstract": "{\\normalsize Dominant processes of neutrino production and neutrino-induced \\ep-pair production are examined in the model of a disk hyper-accreting onto a Kerr black hole. The efficiency of plasma production by a neutrino flux from the disk, obtained for the both cases of presence and absence of a magnetic field, is found to be no more than several tenths of percent and, therefore, not enough for the origin of cosmological gamma-ray bursts.} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0201/hep-th0201101.txt": {
        "abstract": "We consider the equilibria of point particles under the action of two body central forces in which there are both repulsive and attractive interactions, often known as central configurations, with diverse applications in physics, in particular as homothetic time-dependent solutions to Newton's equations of motion and as stationary states in the One Component Plasma model. Concentrating mainly on the case of an inverse square law balanced by a linear force, we compute numerically equilibria and their statistical  properties. When all the masses (or charges) of the particles are equal, for small numbers of points they are regular convex deltahedra, which on increasing the number of points give way to a multi-shell structure. In the limit of a large number of points we argue using an analytic model that they form a  homogeneous spherical distribution of points, whose spatial distribution appears, from our preliminary investigation, to be similar to that of a Bernal hard-sphere liquid. ",
        "introduction": "\\news  This is one of a series of papers about central configurations and related problems involving the equilibria of point particles under the action of two-body central forces.  The main point of the present work is to survey what is known mathematically from a wide range of disciplines and to link this together with some new, mainly numerical, results of our own, establishing a basis for future work on the subject.  Our main emphasis here will be on the classical problem of finding central configurations  of particles associated with an inverse square  interaction force which are trapped by a linear force, induced by a harmonic potential.  Such models are very common in a wide variety of physical applications, but most of our discussion will focus on systems of gravitating points which in addition to the usual attractive inverse square force, experience a repulsive force proportional to their distance from the origin. They arise naturally when seeking homothetic time-dependent solutions of Newton's equations of motion for gravitating point particles, which in turn may have some relevance to Newtonian Cosmology and  models for the large-scale structure of the universe.  Another physical interpretation arises when the inverse square force is thought of as an electrostatic repulsion and the linear force as an attraction,  due to a uniform background of the opposite charge. In this guise the problem originally arose in J.J.  Thomson's static Plum Pudding model of the atom~\\cite{JJ1} in which the positive electric charge is smeared out into a uniform ball (the pudding) while the negatively charged electrons correspond to the plums. Although Rutherford's experiments conclusively  demonstrated that this model is not relevant as a theory of atomic structure, it nevertheless continues to offer insights into the structure of metals (with the role of positive and negative charges interchanged) and other condensed matter systems and is often referred to as the One Component Plasma (OCP) model~\\cite{Bau}, or sometimes as classical Jellium.  Central configurations  are the critical points of a suitable potential function and those configurations which minimize it are numerically the easiest to study. In fact almost all of this paper will be concerned with central configurations which are local minima that coincide with, or are very close to, the absolute minimum of the potential; only in the case of small numbers of points $(\\le 100)$ will we claim to have found the absolute minima. We use two different numerical techniques to compute these minima.  Firstly, a simple multi-start gradient flow algorithm which, given a set of random initial conditions, finds the path of steepest descent toward a local minimum. The other technique is that of simulated annealing~\\cite{sa}, which uses thermal noise to deter the system from falling into a local minimum which is not the global one. We used these two methods in tandem to increase our confidence in finding the true minimum for small numbers of points and to find a stationary point close to the true minimum for larger numbers. By running the codes many times when the number of points is large, we were able to deduce that there are very many local minima with energies close to the absolute minimum. In this regard it resembles related problems such as that of placing point charges on a sphere and those of sphere packing.  It turns out that the Plum Pudding interpretation provides the key to understanding the properties of central configurations for moderate and large numbers of points and is also quite valuable for understanding the solutions for small numbers of points. The idea is that for many purposes one may envisage the equilibria as a packing of Thomson-type hydrogen atoms, that is, electrically neutral spheres containing a single negative charge at the centre in a shell of positive charge.  More quantitatively the spheres correspond to the Thomson atoms described above. The fact that this correspondence may be elevated to a precise quantitative tool was apparently first recognized by Leib and Narnhoffer~\\cite{Lieb} who used it to obtain a rigorous lower bound for the energy of the OCP in terms of a close packing of Thomson atoms. Our numerical results show that the actual minimum is incredibly close to the Leib-Narnhoffer bound and leads to a picture of the equilibria not unlike Bernal's random close packing model of liquids~\\cite{bernal}. We use the word liquid deliberately because despite the wide-spread belief that in the limit of infinite numbers of particles the minimum of the OCP model is given by a Body Centred Cubic (BCC) crystal, our preliminary results for up to 10,000 particles appear to show no sign of crystallization, nor long range translational order. They are, however, crudely consistent with a Bernal liquid.  A second piece of intuition which appears to be useful  is to consider points uniformly distributed inside a sphere. Remarkably, by using the continuum limit, an analytic expression can be derived for the probability distribution for separations in terms of the radius of the confining sphere, which is known in terms of the number of points. This two-point function provides an analytic test of the homogeneity of the distribution, which is passed with considerable accuracy. It is also possible to compute a three-point statistic associated with the distribution of triangles, and we find agreement there too.  We will present our results for various values of the number of points, $N$, in three groups designed to exemplify the specific characteristics of the solutions:  \\begin{itemize} \\item [(I)] Small numbers of points, $N\\le 100$ say.  \\item [(II)] Moderate numbers of points, say  $100 < N < 1000.$  \\item [(III)] Large numbers of points. Here we are able to deal with $1000 \\le N \\le 10,000$. \\end{itemize}  \\noindent For the most part we will stick to the case where all the masses (charges) of the particles are equal ($m_1=m_2= .. =m_N=m$).  \\medskip \\noindent A summary of the results is as follows : \\medskip  \\noindent In case (I) we claim to have found the absolute minima by using the two different algorithms with a wide range of different initial conditions. For $N\\le 12$ the points lie at the vertices of a polyhedron which is a deltahedron (one made entirely from triangles) except for the antiprism found for $N=8$, and is regular if $N=4,6$ or $12$. The polyhedron is a tetrahedron if $N=4$, an octahedron if $N=6$ and an icosohedron if $N=12$. When $N=13$ the minimum is a single point surrounded by the other twelve in an icosahedral structure and for $13\\le N\\le 57$ and $N=60$ there are effectively two shells. There is a link between $N=13$ being the first value at which a point is found inside the polyhedron and the fact that at most 12 spheres of equal radius can touch a given sphere of the same radius. For $58\\le N\\le 100$ (except $N=60$ which is a particularly symmetric structure) there are three shells.  \\medskip \\noindent In case (II) the configurations found by our algorithms, which are local minima but may not be the absolute minima, look at first glance to be roughly uniform. However, closer examination of the precise distribution of points reveals a clearly defined system of shells. For example, if one plots the density as a function of radius it oscillates around uniformity with a regular period. Each of the shells appears to have roughly the same surface density and the radii of the shells appear to be in arithmetic progression. This leads to an approximate description of the number of points in each shell. As the number of points increases the minimum of the energy comes closer and closer to the lower bound, suggesting that the assumptions under which it is derived provide a good picture of the distribution of the particles.  \\medskip \\noindent In case (III) we see that a clear spatial uniformity of the distribution emerges. This is exemplified by computing two-point and three-point statistics and comparing them to the continuum description of the problem. With a few minor caveats related to the discreteness of the distribution, we find remarkable agreement between the analytic expressions and those found for large $N$; the results for the values $N=1000$ and $N=10,000$ will be presented. This uniformity of the density distribution is a  consequence of Newton's theorem: for an inverse square law, the force due to a spherically symmetric  distribution of matter is the same as if the total mass is concentrated at the centre of mass. This is not the case for any other force law. Of considerable interest is the spatial distribution of the particles in these uniform distributions. We computed the distribution of the distance between nearest neighbours and found it to be sharply peaked, suggesting that each particle can be thought of as a sphere of fixed radius and that they may pack as in the classical sphere packing problem. However, a preliminary investigation of the angular distribution of nearest neighbours reveals no evidence of long-range orientational order as one might expect, for example, in a solid. The main caveat to this result is that for large values of $N$ we are unable to have much confidence in having found the global minimum of the energy. Nonetheless, the asymptotic approach to the lower bound on the energy suggests that the configurations we have found are very close to the global minimum. ",
        "conclusions": "\\news \\label{sec-conc}  By use of numerical algorithms we have investigated in detail  central configurations where the interaction force is that of an inverse square law and the masses (charges) of  all the particles are equal. We find that for low values of $N$ the configurations are generally convex deltrahedra which gives way to a multi-shell structure for $N>12$. As $N$ increases the number of shells increases and eventually the configuration tends towards having a constant density. The two-point probability distribution and also the probability of acute angle triangles agree to a high degree with those of a uniform distribution. The distribution of nearest neighbours is sharply peaked suggesting that each particle can be approximated by a sphere of diameter $d\\approx 1.65$ and we have found, at this stage, no evidence for long-range orientational order in contrast to the situation in 2-dimensions, which we shall present elsewhere. It still remains an open question as to whether crystallization occurs, and the possibility remains that for large values of $N$  either we may not have found a minimum sufficiently close to the global one, or that we have not probed sufficiently large values of $N$. These aspects are currently under further investigation.  \\begin{figure} \\begin{center} \\leavevmode \\centerline{\\epsfxsize=10cm\\epsffile{dendif.ps}} \\caption{The density as a function of radial distance in a model where the interaction is generated by an inverse cube force with $N=10000$. Notice that the density is not constant on the outer extremities.} \\label{dendif} \\end{center} \\end{figure}  The specific types of central configurations that we have computed are examples for just one of a large set of models. As we have explained, the interaction potential we have studied has a number of special properties, and we should note that different force laws will lead to very different results. To illustrate this we have included fig.~\\ref{dendif} which shows the density distribution as a function of radial distance for particles with $N=10000$ when the interaction force is an inverse cube law. Clearly in this case there is a decreasing trend in the density with increasing radius, rather than the approach to uniform density that we have encountered so far in this paper. If the power in the interaction force is further increased then this downward trend becomes even more apparent.  \\begin{figure} \\begin{center} \\leavevmode \\ \\vskip -0cm \\centerline{\\epsfxsize=12cm\\epsffile{fig18_bw.ps}} \\caption{The distribution of 124 points, 100 of which have $m=1$, 20 have $m=5$ and 4 have $m=25$. Each point with $m=1$ is represented by a sphere of diameter $d=1.65$ and the others by spheres with a diameter related to their mass by $d\\propto m^{1/3}$. \\goodbreak {\\em See fig18.jpg for a colour version of this figure.} } \\label{diffmass} \\end{center} \\end{figure}  Another interesting possibility is to consider situations in which the particles have different masses. Using the intuition that each of the particles can be represented by a sphere, our earlier analysis suggests that the diameter of this sphere should be taken to be proportional to $m^{1/3}$ and indeed we find this to be the case. This is illustrated in fig.~\\ref{diffmass}, where the spheres can be seen to fit snugly together using the above prescription of taking the volume of the sphere proportional to the mass of the particle.  Clearly the current work is only the tip of the iceberg in terms of the full generality of the concepts involved in central configurations, but we believe it represents a good starting point for further work. Investigations  into different power laws for the interactions,  different mass distributions and the all important question of whether crystallization occurs in these types of models are all underway.\\\\"
    },
    "0201/astro-ph0201352_arXiv.txt": {
        "abstract": "We study a large set of high spatial resolution optical rotation curves of galaxies with the goal of determining the model parameters for a disk embedded within a cold dark matter (CDM) halo that we model either with a Navarro, Frenk \\& White (NFW) profile or pseudo-isothermal profile. We show that parameter degeneracies present in lower resolution data are lifted at these higher resolutions. 34\\% of the galaxies do not have a meaningful fit when using the NFW profile and 32\\% when using the pseudoisothermal profile, however only 14\\% do not have a meaningful fit in either model.  In both models we find correlations between the disk baryon fraction $f_d$ and the spin parameter of the halo $\\lambda'$, between $f_d$ and the dark halo mass $M_{200}$, and between $M_{200}$ and the concentration parameter $c$. We show that the distribution of the concentration parameters $c$, for a NFW halo, is in good agreement with CDM predictions; no significant galaxy population is found with very low values of $c$. The overall distribution of $\\lambda'$ is in good agreement with theoretical predictions from hierarchical tidal torque theory. The whole sample is also well fitted by a pseudo-isothermal dark halo with a core, but the size of the core is rather small (6\\% of the virial radius or smaller; for 70\\% of the sample the core size is less than 2 kpc). Thus we conclude that the profile of dark matter is steep ($r^{-1}$ or steeper) down to this radius; large dark matter cores (and therefore very low dark matter central densities) seem to be excluded. LSBs tend to have higher values of $\\lambda'$ for a given $f_d$ and lower values of $c$ for a given mass than HSBs.  In an appendix we give some useful formula for pseudo-isothermal profile halos and discuss in detail the issue of parameter degeneracies. ",
        "introduction": "The Cold Dark Matter (CDM) paradigm for structure formation has proved remarkably successful in explaining the observed large-scale properties of the universe, such as the abundance and clustering of galaxies and clusters \\citep{Peacock+01,Verde+2dF01,Lahav+2dF01}, the statistical properties of the Ly$\\alpha$ forest (e.g., \\citet{Croft+01}), and the power spectrum of the cosmic microwave background anisotropies (e.g., \\citet{Jaffe+01}). However, a number of puzzling discrepancies remain when CDM predictions are extrapolated to small scales. Chief among them are: (i) The substructure problem. CDM over-predicts the number of satellites around a Milky-Way sized galaxy by an order of magnitude \\citep{KKVP99,MGGLQST99}. (ii) The density profile problem. Observed rotation-curves of dwarf and lower surface brightness galaxies suggest that the inner regions have a constant density core rather than the density cusp predicted by CDM (e.g., \\citet{MB98,MGGLQST99,DB00}).  Furthermore, numerical SPH simulations of disk formation within dark halos fail to  match the zero-point of the observed Tully-Fisher. One possible solution is to lower the central concentration of the dark matter halos \\citep{NS00}.  These problems have stimulated numerous proposed astrophysical solutions, as well as modifications of the fundamental CDM paradigm itself (e.g., \\citet*{SS00,KL00,CDW96,BOT01,Goodman00,HBG00,Cen01}).  In this paper we once again examine the density profile problem by fitting CDM models to observed rotation curves of spiral galaxies.  Our study differs from previous work in two respects. Firstly, we use optical rotation curves rather than the HI rotation curves used in previous studies which are plagued by beam smearing \\citep{SMT00,vanBRDB00}. Even if the effects of beam smearing are neglected, the relatively large errors and limited spatial sampling of HI rotation curves imply that they cannot be used to discriminate between constant density cores and $r^{-1}$ cusps \\citep*{vanBS01}. By contrast, the optical rotation curves we use are free from beam smearing, have smaller errors and higher spatial resolution. We show that, with this superior data, various parameter degeneracies present when fitting HI rotation curves can be lifted.  This allows us to distinguish between core and cusp-like inner profiles.  Secondly, while many studies have focused on relatively small samples of dwarfs (but see \\citet{Navarro98}), we use a large sample (400 galaxies) spanning a wide range in luminosity and surface brightness.  Our goal is to determine the best fitting model parameters for a disk within the CDM profile proposed by Navarro Frenk and White (1997; NFW) and within a pseudo-isothermal profile halo and study the distributions of the recovered disk parameters.  We find that the NFW profile provides a good fit to 66\\% of the galaxies in the sample, with a distribution of recovered concentration parameters and spin parameters broadly consistent with that predicted by CDM numerical simulations. When the sample is fitted with an isothermal profile with a core, 68\\% of the galaxies are well fitted whithin this model and the best fit parameters favor cores with small sizes (below 6\\% of the dark halo virial radius).  Many rotation curves that have no meaningful fit in one dark matter profile are well fitted by the other. There is no meaningful fit in either model for only for 14\\% of the sample. This is consistent with the inner dark matter profiles being steep (slope of -1 or steeper) down to radii that are few percent of the virial radius.  In both models the recovered baryonic mass to light ratios are broadly in agreement with predictions from synthetic stellar populations.  This paper is organized as follows.  In section 2 we present the set of rotation curves we analyze and the two models we use to fit the data (exponential disk embedded in a dark matter halo with a NFW profile or a pseudo-isothermal profile). These two models have four free parameters, but since for most of the galaxies the scale length of the disk is given, we report the analysis for four free parameters in appendix II and present in the main text the results for the fit with three free parameters.  In section 3 we illustrate and discuss the correlations we find between the best fit parameters for the two models. We discuss the general implications of these findings in section 4 where we also compare the derived baryonic mass-to-light ratios with the range allowed by stellar populations.  Finally, we conclude and summarize our results in section 5.  In appendix I we develop the expression for the rotation curve of an exponential disk embedded in a pseudoisothermal halo. In appendix III we study the degeneracies among the parameters of the models. ",
        "conclusions": "We have analyzed a large set of high spatial resolution rotation curves of galaxies with the goal of determining the model parameters for a disk embedded within a cold dark matter (CDM) halo that we have modeled in two ways: either with a NFW profile or a pseudo-isothermal profile.  To this aim we have developed the expression for the rotation curve of an exponential disk embedded in a pseudoisothermal halo (Appendix I).  Our study differs from previous work in two respects. Firstly, we use optical rotation curves rather than HI rotation curves, so the rotation curves we use are free from beam smearing and have smaller errors and higher spatial resolution.  Secondly, while many studies have focused on relatively small samples of dwarfs, we use a large sample (400 galaxies) spanning a wide range in luminosity and surface brightness.  We find that the NFW profile provides a good fit to 66\\% of the galaxies in the sample (in agreement with previous studies, \\citet*{vanBS01}) that were based on few (20) dwarf galaxies), with a distribution of recovered concentration parameters and spin parameters broadly consistent with that predicted by CDM numerical simulations. When the sample is fitted with an isothermal profile with a core, 68\\% of the galaxies are well fitted within this model and the best fit model favors cores with small sizes (almost all have $r_{c} < 0.06 R_{200}$ and $\\sim 70\\%$ have $r_{c} < 0.01 R_{200}$). However, we find that many rotation curves that have no meaningful fit for one dark matter profile are well fitted by the other: for only 52 galaxies (14\\% of the sample) there is no meaningful fit in either model. These largely comprise low quality rotation curves. Our findings are consistent with the inner dark matter profiles being steep (slope of -1 or steeper) down to radii that are few percent of the virial radius. Large dark matter cores (and therefore low central dark matter densities) seem thus to be excluded.  Numerical SPH simulations of disk formation within dark halos find that disks loose a significant fraction of their angular momentum if: 1) the dark matter profile is steep in the center and 2) the baryons settle in right after virialization \\citep{NSbis00}. On the other hand, we recover steep dark matter profiles and no significant loss of angular momentum. There are several routes to explain this.  Some authors have pointed out that if baryons are allowed to settle well after virialization of the halo (e.g. \\citet{SDTC01}), then there is no significant loss of angular momentum. Another possibility to consider is that current hydrodynamical simulations lack the resolution to simulate a multi-phase interstellar medium and also to follow the detailed formation of giant molecular clouds, and therefore the sites of star formation, within the settling disk. It is therefore not clear that current simulations contain all the physical ingredients needed to simulate the complex process of baryons cooling down into a disk.  In both models, for the galaxies that have meaningful fit to the rotation curves, the recovered baryonic mass to light ratios are broadly in agreement with predictions from synthetic stellar populations.  From this large sample of galaxies we find several correlations among the best fitting model parameters and some regions in the parameters space that are not populated, we refer to those regions as zone of avoidance.  These correlations an zone of avoidance are remarkably similar in both profiles and do not change when the rotation curves are fitted by a model with three free parameters or with four free parameters.  In particular we find a strong correlation between the disk mass fraction and the spin parameter for different kind of galaxies (HSBs and LSBs), in good agreement with findings of previous studies that involved only few dwarf LSBs galaxies. These correlations may be a result of parameter degeneracies if errors in the rotation curve have been underestimated; otherwise, they could have a physical origin.  When fitting the sample with the NFW profile we obtain that for the 66\\% of galaxies for which meaningful fits to the rotation curves are found, the distribution of concentrartion parameters is in good agreement with predictions from N-body simulations. Also, the distribution of disk spin parameters is in broad agreement with the distribution of spin parameters for the dark matter predicted from N-body simulations and linear theory, but with a small tail at low spins.  The primary motivation in previous work of studying the rotation curves of dwarfs and/or LSB's was that they would be dark matter dominated. Although it is quite plausible, this claim must be taken with caution.  In fact, our analysis of HSB and more luminous galaxies, as we have seen, yields values for the baryonic disk mass fraction very similar to previous analysis of dwarfs; a naive interpretation of rotation curve fits alone would lead us to claim these galaxies as dark-matter dominated as well. In Appendix II, we analyze the LSB sample of \\citet*{MRB01} and find that a somewhat larger fraction of them cannot be fit by NFW profiles, $\\sim 57 \\%$ as opposed to $\\sim 33\\%$ for the entire sample (those which do fit have somewhat low but theoretically plausible concentration parameters $c$). This difference may be a 1 $\\sigma$ statistical fluctuation (there are only 26 LSBs in the sample), it may reflect the greater fidelity of LSB rotation curves to the true underlying dark matter distribution, or it may perhaps be due to a systematic difference in the dark matter profile of halos which undergo greater tidal torquing. On the other hand, for the pseudoisothermal model, only one LSB has $f_d > 0.2$. This seems to indicate that LSBs are indeed fitted better with a dark matter core. For the LSB sample, 9 galaxies have dark matter cores with sizes below 1.5 kpc and 16 below 3 kpc. This is in fair agreement with our previous findings that the size of the dark matter core is modest. It is also in agreement with the dark matter core size inferred from inverting the Poisson equation using the observed rotation curve \\citep{BMBR01}. Ultimately, our sample choice involves a trade-off: we gain increased precision in the rotation curve (which breaks parameter degeneracies) and a much larger sample encompassing a much broader class of galaxies, but also acquire a somewhat larger uncertainty in the disk contribution to the rotation curve.  In general, LSB galaxies tend to have higher values of $\\lambda'$ for a given $f_d$ and lower values of $c$ for a given mass than HSB galaxies (see Appendix II). This supports the view that the dark matter profile of LSBs might reflect an enviroment in which dark halos undergo a higher tidal torque.  If the $M/L$ of the disk was well constrained, this ambiguity can be eliminated: it would be possible to subtract the disk contribution to the rotation curve and recover directly the inner profile of the dark halo by inverting the Poisson equation (e.g., \\citet{BMBR01}). The M/L can be constrained from multi-color photometry or even better, high $S/N$ spectra of the stellar population in galaxies (extending as close as possible to the $K$ band) together with high-spatial resolution rotation curves especially of the inner 1-2 Kpc. A more detailed treatment of this issue will be presented in a forthcoming paper."
    },
    "0201/astro-ph0201164_arXiv.txt": {
        "abstract": " ",
        "introduction": "Supernovae observations have recently provided evidence that the expansion of the Universe is undergoing a late time acceleration \\cite{cc,SCP,Riess:2001gk}.  This acceleration can be explained in the framework of standard cosmology by a non vanishing cosmological constant. Although in agreement with current observations, such an explanation exacerbates the usual cosmological constant problem because it requires an explanation for its  very small, but non zero, value.  One may wish to find alternative explanations for the acceleration, and there are several proposals in the literature. Here we explore a scenario proposed  in \\cite{Deffayet:2001uy,Fifth},  based on the model of Dvali-Gabadadze-Porrati of brane-induced gravity \\cite{DGP}. This proposal explains the observed late time acceleration of the expansion of the Universe through a large scale modification of gravity coming from ``leakage'' of gravity at large scale into an extra-dimension, and without requiring a non vanishing cosmological constant. The interesting point about this model from a phenomenological perspective is that it is a testable alternative to a cosmological constant model with the same number of parameters. This is in contrast to models of ``quintessence'' where the equation of state of the new component becomes a free function that needs to be constrained.  In \\cite{Fifth} it has been shown that the model was in qualitative agreement with all known cosmological observations. The purpose of this work is to go one step further and quantitatively confront the model with observations of supernovae and the cosmic microwave background (CMB).  The paper is organized as follows, in section \\ref{Gravity} we discuss the dynamics of the background metric of the universe in the model. We first introduce in a few words the brane-induced gravity model of Dvali-Gabadadze-Porrati  \\cite{DGP} (see also \\cite{DG,Dvali:2001gm,Dvali:2001gx}) which provides the framework (subsection \\ref{branind}). We then discuss the cosmological dynamics for the accelerated solution considered in this paper (subsection \\ref{cosmdyn}). In the following, we confront the model with the Supernovae observations of the Supernova Cosmology Project (SCP) \\cite{SCP} (subsection  \\ref{SNsect}) and CMB data (subsection  \\ref{CMB}).  Our fits indicate that the model is currently in agreement with SNIa and small scale CMB data. One can hope to discriminate the model from standard cosmology using future precision cosmological parameters measurements, but also maybe  modifications in the growth of large scale structures. ",
        "conclusions": "The fits done in this work show that the model of accelerated universe through gravitational leakage into extra dimension of Ref. \\cite{Deffayet:2001uy,Fifth} is in current agreement with SNIa and CMB data.  The degeneracies in parameters estimations using one data set (e.g. CMB) can be partially lifted using the other (e.g. SNIa) as in standard cosmology. The Supernovae data prefer a slightly lower value of $\\Omega_M$ ($\\Omega_M = 0.18^{+0.07}_{-0.06}$) than the CMB for a flat universe, however a concordance model with $(\\Ok, \\O5, \\od, $ $ \\ob, n, A) = (0,0.1225,0.1,0.02,0.96,0.57)$ which has $\\Om=0.3$ (and $\\chi^2\\approx140$ for the full data set (135 data points)) provide a good fit to both sets, all the more as we have not included systematic errors in our parameter estimations.  For this model the crossover distance between 4D and 5D gravity is given by $r_c \\sim 1.4~H_0^{-1}$.    We have also given the equation of evolution for cosmological perturbations.  Those equations were used to justify the approximation we made to compute cosmological perturbations, namely we used standard four dimensional evolutions equations over a background with a scale factor given by the accelerated solution given in \\cite{Fifth}. This is justified for small scale CMB anistotropies (scale smaller than the crossover scale $r_c$). From those equations, and the known behavior of gravity in the model at hand, one can also expect modifications in the growth of large scale structure. This could potentially lead to a way to discriminate between standard cosmology and the model considered in this work, and is left for future investigation.   We want to end by noting that the model under consideration is very predictive in the sense that future observations have the potential to rule it out. In contrast to quintessence models, this model has the same number of free parameters as the usual LCDM model. With the advent of new precision cosmological measurements such as new SNIa observations, CMB measurements, ongoing galaxy surveys such as Sloan and 2dF, weak lensing surveys, etc. it should be possible to test the model very accurately (for a recent summary of how different observations will constrain the matter content of the universe see \\cite{Tegmark:met} and references therein)."
    },
    "0201/astro-ph0201487_arXiv.txt": {
        "abstract": "We present spectroscopy and time-series photometry of the newly discovered dwarf nova 1RXS J232953.9+062814.  Photometry in superoutburst reveals a superhump with a period of 66.06(6) minutes.  The low state spectrum shows Balmer and HeI emission on a blue continuum, and in addition shows a rich absorption spectrum of type K4 $\\pm$ 2. The absorption velocity is modulated sinusoidally at $P_{\\rm orb} = 64.176(5)$ min, with semi-amplitude $K = 348(4)$ km s$^{-1}$.  The low-state light curve is double-humped at this period, and phased as expected for ellipsoidal variations.  The absorption strength does not vary appreciably around the orbit. The orbital period is shorter than any other cataclysmic variable save for a handful of helium-star systems and V485 Centauri (59 minutes). The secondary is much hotter than main sequence stars of similar mass, but is well-matched by helium-enriched models, indicating that the secondary evolved from a more massive progenitor.  A preliminary calculation in which a 1.2 M$_{\\odot}$ star begins mass transfer near the end of H burning matches this system's characteristics remarkably well. ",
        "introduction": "Cataclysmic variable stars (CVs) are close binary systems in which a low-mass secondary transfers mass onto a white dwarf; \\citet{warn} wrote an excellent monograph on CVs. Orbital angular momentum losses $\\dot J$ evidently drive CV evolution.  As the orbit shrinks, the secondary star's Roche critical lobe contracts, causing mass transfer. The Roche geometry tightly constrains the secondary star's mass at a given orbital period $P_{\\rm orb}$.  Short-period systems have low-mass secondaries, so if the chemical composition is normal ($X \\sim 0.7$), the secondary is faint and contributes negligibly to the visible-light spectrum (Fig.~4 of \\citealt{patprecess01}). For normal compositions the radius is expected to reach a minimum around 0.05 $M_{\\odot}$, leading to a predicted period minimum around 70 -- 75 minutes (the exact value being dependent on $\\dot J$), with subsequent evolution driving the system to greater separations \\citep{kb99,patlate98,ps81}.  During outburst, short-period dwarf nova systems often show photometric oscillations (superhumps) at periods a few percent longer than $P_{\\rm orb}$. The superhump frequency is thought to be the beat between $P_{\\rm orb}$ and a tidally driven precession of an eccentric disk.  The fractional period excess of the superhump appears to be a measure of the mass ratio \\citep{patprecess01}.  We recently observed 1RXS J232953.9+062814 (hereafter RX2329+06), a newly-recognized dwarf nova system with $P_{\\rm orb}$ below the canonical minimum period for hydrogen-rich secondary stars. Our observations reveal an unexpectedly hot secondary star. We suggest that the secondary has undergone substantial nuclear evolution leading to an enhanced helium abundance. ",
        "conclusions": "It is surprising to find such a hot secondary at this short period, where the secondary mass must be $\\sim 0.1$ M$_{\\odot}$. No calculation beginning with a solar-abundance secondary has produced periods this short \\citep{kb99}, and hydrogen-rich stars of this small a mass would be much cooler than observed.  We suggest that the secondary is substantially enriched in helium, as a result of nuclear evolution prior to mass transfer. As shown by \\citet{bk00}, secondary donors which have evolved off the Zero Age Main Sequence (ZAMS) at the onset of mass transfer can explain a substantial fraction of the observed CVs with late spectral types and $P_{\\rm orb} > 6$ h. Test calculations with constant mass transfer rates show that significantly evolved donors (e.g. with mass transfer starting near the end of central H burning) can be much hotter at a given $P_{\\rm orb} \\simle 5$ h than ZAMS donors. This is illustrated in Fig. \\ref{fig4}, which displays sequences with initial donor mass $M_2$ = 1.2 $\\msol$ and starting mass transfer on the ZAMS (solid line) or near the end of H burning (dashed line). As already noted by \\citet{bk00}, such extreme sequences never become fully convective, because of the lower central H abundance, and may continue to transfer mass in the 2-3 h period gap. If they sequences do so, they can reach very low orbital periods with unusually high $\\te$.  For the test case shown in Fig.~\\ref{fig4}, the evolved sequence (dashed line) reaches $P_{\\rm orb}$ = 64 min with a mass $M_2 \\sim$ 0.11 $\\msol$ and a radius $R_2$ $\\sim$ 0.13 $R_\\odot$, in close agreement with the estimates in \\S 3. We derive a spectral type K5, based on the empirical SpT - $(I-K)$ relation of Beuermann et al (1998), again matching observation. We stress that our spectral type estimate is to be treated with caution, since the surface chemical composition is expected to be non-standard, with a mass fraction of H  $\\sim$ 30-40\\%. Because CNO-processed material should be visible on the surface, the model has C significantly depleted and N enriched by a factor $\\sim$ 5, whereas O is hardly affected.  We stress that this evolutionary sequence is preliminary, but it provides a surprisingly good match to the observed properties of RX2329.  The emission line fluxes (Table 1) suggest that helium is enhanced. The ratio of H$\\alpha$ to HeI $\\lambda6678$ is about 3.6.  We measured the H$\\alpha$/$\\lambda$6678 ratio in archival spectra of SU UMa stars \\citep{tpst,tt97,thor97} and found typical values of 8, with none below 6.  This could be an excitation effect, but even high-excitation novalikes such as V603 Aql \\citep{patt97} have larger ratios.  The 59-minute binary V485 Cen also appears to have an unusually low H$\\alpha$/$\\lambda$6678 in the spectrum published by \\citet{aug96}, but absorption features are not evident there. As noted earlier, the relative CNO abundances should be affected, but we cannot comment on this since CNO elements lack strong lines at this $\\te$.  In sum, RX2329 is a CV in which the secondary evidently has undergone significant nuclear burning and then had much of its mass stripped away, since at its present mass, its nuclear evolution timescale greatly exceeds the Hubble time. If the white dwarf mass is similar to most CVs ($\\sim 0.7$ M$_{\\odot}$), much of the mass lost by the secondary appears to have been lost from the system.  Evidently the system formed with a secondary considerably more massive than the white dwarf.  In principle this leads to mass transfer on a thermal timescale, until $M_2$ becomes small enough for the system to reappear as a standard CV (see \\citealt{bk00} for a discussion).  {\\it Acknowledgments.} We gratefully acknowledge funding by the NSF (AST 9987334), and we thank the MDM staff for their excellent support.  Special thanks go to the CBA observers who contributed to the photometry, including Dave Skillman, Arto Oksanen, Ed Beshore, and Tonny Vanmunster; a fuller report on their work will be forthcoming.  We made use of the USNOFS Image and Catalogue Archive operated by the United States Naval Observatory, Flagstaff Station (http://www.nofs.navy.mil/data/fchpix/).  {\\it Note added 2002 January 29:  We obtained seven more spectra with the 2.4m on 2002 Jan. 21 UT.  Combining absorption velocities from these with the 2001 November data yields a refined $P_{\\rm orb}$ = 0.0445671(2) d.}  \\clearpage"
    },
    "0201/astro-ph0201214_arXiv.txt": {
        "abstract": "Methods used in the radial-velocity program of short-period binary systems at the David Dunlap Observatory are described with particular stress on the Broadening Function (BF) formalism. This formalism has permitted determination of radial velocities from complex spectra of multiple-component systems with component stars showing very different degree of rotational line broadening. The statistics of random errors of orbital parameters is discussed on the basis of the available orbital solutions presented in the six previous papers of the series, each with ten orbits. The difficult matter of systematic uncertainties in orbital parameters is illustrated for one typical case of GM~Dra from the most recent Paper~VI. ",
        "introduction": "\\label{sec1}  This paper should be considered as a companion and supplement to the previous papers of our series of radial velocity studies of close binary stars: \\citet[Paper I]{ddo1}, \\citet[Paper II]{ddo2}, \\citet[Paper III]{ddo3}, \\citet[Paper IV]{ddo4}, \\citet[Paper V]{ddo5}, \\citet[Paper VI]{ddo6}.  The current program of radial velocity observations of close binary systems with periods shorter than one day is approximately at its half-way point. Our methods have been evolving slightly during the execution of the 60 radial velocity orbits presented in the six papers of the series, but appear to have stabilized now, warranting a more detailed documentation of the essential steps in our analysis and data reductions. We summarize these methods and give an overview of the uncertainties so that the results described in the previous and the planned future papers of the series could be better evaluated by readers. The discussion is limited strictly to methodological aspects and does not include any astrophysical results which will be discussed after the program is concluded. ",
        "conclusions": "\\label{plans}  The ongoing survey of close binary systems with periods shorter than one day, currently conducted at the David Dunlap Observatory, has resulted in a consistent set of radial velocity orbits for sixty previously unobserved binaries to approximately 11th magnitude. While, at the start, the survey concentrated on systems which simply had not been studied before (for various reasons, but mostly because of inadequate instrumentation and data-analysis tools some half a century ago, when this field was very active), the photometric discoveries of the Hipparcos satellite are now dominating in numbers. There was only one Hipparcos system among the first twenty orbits (Papers I and II), 9 such systems among the next twenty orbits (Papers III and IV) and 15 such systems among the most recent twenty orbits (Papers V and VI). About 50 known, photometrically-discovered binaries still remain to be observed and analyzed and new ones are constantly added to catalogs, some of them quite bright. Regrettably, apparently there is no similar survey for the southern hemisphere.  Our survey is quasi-random in the sense that we observe all short-period ($P < 1$ day), bright, previously unobserved binaries. With such criteria, the contact binaries absolutely dominate in numbers. Among the 60 systems described in the previous six papers, only 8 were not contact systems. This is partially due to strong selection effects against detection of detached binaries, but mostly due to the very high frequency of contact binary systems in the old-disk population, particularly in the period range 0.3 to 0.5 days, but with a tail extending beyond one day, to about 1.3 -- 1.5 days. The high frequency of incidence is strongly manifested in the volume-limited OGLE sample and in open clusters \\citep{ruc98}. Because our survey is magnitude limited, we tend to include many brighter systems from the tail of the distribution between 0.5 day and our current upper limit at one day. Otherwise, we do not discriminate among binary systems in any other way. In particular, the random character of the survey has resulted in discoveries of the largest ($q=0.97$, V753~Mon; Paper~III) and the smallest ($q=0.066$, SX~Crv; Paper~V) known mass ratios among contact binaries.  The DDO survey is characterized by moderate random errors of about 1 -- 2 km~s$^{-1}$ for the orbital parameters, $V_0$, $K_1$ and $K_2$, and -- upon completion -- can serve as a useful database of parameters of very close binary systems. We are aware, however, that our final parameters contain systematic uncertainties resulting from our radial-velocity measurement techniques. While the use of the broadening functions permitted us to analyze close binaries in several multiple, visual/spectroscopic systems providing data which were too ``difficult'' before, our extraction of individual radial velocities from the broadening functions, through Gaussian fitting, is a disputable approach for contact binary systems. Because the line-broadening for such systems is very strong, comparable with orbital velocities of hundreds of km~s$^{-1}$, and -- in fact -- somewhat asymmetric, our measuring technique may lead to systematic errors reaching levels of 5 -- 7 km~s$^{-1}$ or even more. Paradoxically, through the use of the broadening functions in place of the cross-correlation functions, we have uncovered real physical reasons why the Gaussian approximation is only barely appropriate. The correct approach avoiding the systematic errors would be to model the broadening functions and determine the radial velocities in terms of the mass ratio, $q$, and the scaling factor, ($K_1+K_2$), with the shift, $V_0$. The models would require independent input from parallel solution of light curves, providing the orbital inclination angle, $i$, as well as the degree-of-contact, $f$. Currently, most of the program targets have not had their light-curves solved, and even if some attempts have been made, we would not trust them for the following simple reason: We have seen so many cases of the spectroscopic mass ratio different from the previous photometric mass-ratio determinations, $q_{sp} \\ne q_{phot}$, that we feel very strongly that the values of $q_{phot}$ are usually not properly constrained and may be plainly wrong\\footnote{Totally eclipsing systems are an exception, as pointed by \\citet{MD72a,MD72b}, but then chances of total eclipses depend on the mass-ratio itself (a wider range of inclinations for small values of $q$), producing a very complex bias in the uncertainties of $q_{phot}$.}, leading to entirely incorrect combinations of orbital parameters.  We envisage that the results of this survey will provide just a first stage of an iterative process. In future, our spectroscopic values of mass ratio, $q_{sp}$, should permit solution of light curves which were previously unsolvable because to the poorly constrained mass ratios. The derived information on ($i$, $f$) pairs would permit, in turn, a re-discussion of the broadening functions and determination of the final orbital parameters, free of systematic uncertainties.  Concerning the instrumental developments at the DDO: Soon, we plan to start using a new CCD system based on a much more sensitive detector. While the analysis of the data should remain the same as described above, we may have to select the targets more discriminately. In particular, it may turn out impractical to observe all binaries with periods shorter than one day down to to the expected limiting magnitude of about 12.5 magnitude. Indeed, from the point of astrophysical usefulness, it would be advantageous to reduce the deficit of the intrinsically faint contact systems among spectroscopically studied binaries of the magnitude-limited sample, by attempting to form a volume-limited sample through giving preference to very short-period systems."
    },
    "0201/astro-ph0201022_arXiv.txt": {
        "abstract": "The dynamical systems of planet-belt interaction are studied by the fixed-point analysis and the bifurcation of solutions on the parameter space is discussed. For most cases, our analytical and numerical results show that the locations of fixed points are determined by the parameters and these fixed points are either structurally stable or unstable. In addition to that, there are two special fixed points: the one on the inner edge of the belt is asymptotically stable and the one on the outer edge of the belt is unstable. This is consistent with the observational picture of Asteroid Belt between the Mars and Jupiter: the Mars is moving stablely close to the inner edge but the Jupiter is quite far from the outer edge. ",
        "introduction": "The discovered number of extra-solar planets is increasing dramatically due to astronomers' observational effort, therefore the dynamical study in this field is getting important. Because the belts of planetesimals often exist among planets within a planetary system as we have in the Solar System, it is indeed important to understand the solutions of dynamical systems of planet-belt interaction. Jiang \\& Ip [2001] predicted that the interaction with the belt or disc might bring the planetary system of upsilon Andromedae to the current orbital configuration.  Yeh \\& Jiang [2001] used phase-plane analysis to study the orbital migration problem of scattered planets. They completely classify the parameter space and solutions and conclude that the eccentricity always increases if the planet, which moves on circular orbit initially, is scattered to migrate outward. These analytical results is consistent with the numerical simulations in Thommes, Duncan \\& Levison [1999].  In addition to astronomy, general or Newton's dynamical systems are studied in many other fields and have very important applications. Clausen et al. [1998] studied periodic modes of motion of a few body system of magnetic holes both experimentally and numerically. Kaulakys et al. [1999] showed that a systems of many bodies moving with friction can experience a transition to chaotic behavior.  On the other hand,  Chan et al. [2001] studied bifurcation for limit cycles of quadratic systems interestingly. Similar type of approach should be also good for the bifurcation of solutions for dynamical systems of planet-belt interaction. In this paper, we focus on the planet-belt interaction and study the bifurcation of such system by phase plane analysis.  Basicly, we would like to understand the orbital evolution of a planet which moves around a central star and  interacts with a belt. The belt is a annulus with inner radius $r_1$ and outer radius $r_2$, where  $r_1$ and $r_2$ are assumed to be constants. We set $r_1=3$ and $r_2=6$ for all numerical results in this paper.  We assume the distance between the central star and the planet is $r$, where $r$ is a function of time. When $r < r_1$, the belt would only give the force which pulls the planet away from the central star. When $r > r_2$, the belt would only give the force which pushes the planet towards the central star. When $r_1 \\le  r \\le r_2$, in addition to the usual gravitational force, there is friction between the planet and the belt. These three cases will be studied in Model A ($r<r_1$), Model B($r>r_2$) and Model C($r_1<r<r_2$) individually.  We will mention our basic governing equations in Section 2. In Section 3-5, we study the locations and stabilities of fixed points by phase plane analysis for Model A, B and C. The conclusions will be in Section 6. ",
        "conclusions": "We have studied the properties of solution curves on the phase plane of dynamical systems of planet-belt interaction by the standard fixed-point analysis.   The system is divided into three regions: (a) the region between the central star and the inner edge of the belt, (b) the region out of the outer edge of the belt, (c) the region between the inner edge and the outer edge of the belt. The dynamics in three regions are governed by three different models and the planet moves around these regions.  We analytically prove some properties of the system and also show some bifurcation diagrams of fixed points for different cases: the number and properties of fixed points changed with the values of parameters $k_2$ and $M_b$ in general.  Our analytical and numerical results show that, in most cases, the locations of fixed points depend on the values of these parameters and these fixed points are either center (structurally stable) or saddle (unstable).  On the other hand,  we found that there are two special fixed points: one is on the inner and another is on the outer edges of the belt. The one on the inner edge is a stable focus (asymptotically  stable) and the one on the outer edge is a saddle (unstable).  This interesting result is consistent with the observational picture of Asteroid Belt between the Mars and Jupiter: the Mars is moving stablely close to the inner edge but the Jupiter is quite far from the outer edge."
    },
    "0201/astro-ph0201508_arXiv.txt": {
        "abstract": "We demonstrate that the spectral differences between Flat Spectrum Radio Quasars (FSRQ) with steep gamma-ray spectra (MeV-blazars) and  FSRQ with flat gamma-ray spectra (GeV-blazars) can be explained by assuming that in the MeV-blazars, the production of gamma-rays is dominated by Comptonization of infrared radiation of hot dust, whereas in the GeV-blazars --- by Comptonization of broad emission lines.  Additional ingredient, required to reach satisfactory unification, is an assumption that the radiating electrons are accelerated via a two step process, in the lower energy range -- following instabilities driven by shock-reflected ions, and in the higher energy range -- via resonant scatterings off Alfven waves. Our model predicts that on average, the MeV-blazars should vary on longer time scales than GeV-blazars. ",
        "introduction": "\\label{mblazejowski-C2_sec:as}  \\noindent  \\begin{itemize} \\item  Relativistic electrons  in a subparsec jet are produced in shocks formed by colliding inhomogeneities. The inhomogeneities are assumed to be intrinsically identical and to move down the jet with bulk Lorentz factors $\\Gamma_2 >\\Gamma_1 \\gg 1$;  \\item Time scale of the collision, as measured in the comoving frame of the shocked plasma, is \\be t_{coll}' \\simeq t_{fl} {\\cal D}  \\, ,\\ee where $t_{fl}$ is the observed time scale of the flare, and \\be {\\cal D} \\equiv {1 \\over \\Gamma (1 - \\beta \\cos{\\theta_{obs}})} \\ee is the Doppler factor of the shocked plasma, and \\be \\Gamma \\simeq \\sqrt {\\Gamma_1 \\Gamma_2} \\, ; \\ee  \\item Inertia of inhomogeneities is dominated  by protons (i.e. $n_e/n_p \\ll m_p/m_e$);  \\item Efficiency of energy dissipation is defined as \\be \\eta_{diss} \\simeq {((\\Gamma_2/\\Gamma_1)^{1/2} - 1)^2 \\over (\\Gamma_2/\\Gamma_1)+1} \\, ,\\ee  \\item Injection of relativistic electrons is approximated by a two-power-law function, with the break at $\\gamma_b$, at which magnetic rigidity of electrons is equal to rigidity of thermal protons, i.e., when their momenta are equal \\be m_e \\sqrt{\\gamma_b^2 -1} \\simeq  m_p \\sqrt {\\gamma_{p,th}^2-1} \\, ,\\ee where \\be \\gamma_{p,th} -1 = \\eta_{p,th} \\kappa  \\ee is the average thermal  proton energy in the shocked plasma, $\\eta_{p,th}$  is the fraction of the dissipated  energy tapped to heat the protons and \\be \\kappa \\simeq { ((\\Gamma_2/\\Gamma_1)^{1/2} - 1)^2 \\over 2(\\Gamma_2/\\Gamma_1)^{1/2}} \\, .  \\ee is the amount of energy dissipated per proton in units of $m_pc^2$. \\noindent \\item  Since the time scales of the flares in FSRQ are rarely shorter than 1 day, the distances of their production, \\be r_{fl} \\sim (r_{fl} /\\Delta r_{coll}) c t_{fl} {\\cal D} \\Gamma \\ee are expected to be larger than 0.1 parsec (where $\\Delta r_{coll}$ is a distance range over which the flare is produced). At such distances, the largest contribution to the energy density of an external radiation field $u_{ext}'$ is provided by the diffuse component of the broad emission lines and infrared radiation of hot dust. \\end{itemize} ",
        "conclusions": ""
    },
    "0201/astro-ph0201058_arXiv.txt": {
        "abstract": " ",
        "introduction": "The majority of the total baryonic matter in the local ($z\\ls 1$) Universe is predicted to be concentrated in highly ionized gas. Structures that are already virialized contain warm ($10^5-10^6$ K) or hot (10$^7$ K) gas [the dense interstellar medium (ISM) of galaxies, and the intracluster medium (ICM) of clusters of galaxies].  The greater amount of baryonic matter is predicted to lie in, as yet unvirialized, matter in the form of a tenuous warm-hot intergalactic medium (WHIM, Hellsten et al., 1998).  The detection and study of these components is needed for the proper understanding of large and small scale structures in the Universe.  However, while X-ray emission from the virialized density peaks of the ICM and ISM has been detected and intensively studied in X-rays, the predicted highly ionized gas in the WHIM has been poorly studied so far, due to instrumental limitations. The low density of the WHIM leads to low emissivity, so that studies of the WHIM in emission are a formidable challenge. However, absorption depends only on the total column density of the medium, not on density, and background light sources in the form of quasars (Aldcroft et al., 1994) and gamma-ray bursts (Fiore et al. 2000) are readily available.  A few high ionization transitions, notably OVI $\\lambda$1031.93, lie in the UV, but the most prominent ions (CVI, OVII, OVIII, NeIX) have their strongest transitions in the soft X-ray band (10-40 \\AA) for a wide range of temperatures ($10^5-10^{6.5}$ K, Nicastro et al., 1999), which should give rise to an ``X-ray Forest'' of absorption lines (e.g. Perna \\& Loeb, 1998, Fang \\& Canizares, 2000).  The advent of high resolution ($R\\sim$1000) soft X-ray spectroscopy with {\\em Chandra} and XMM-{\\em Newton} allows sensitive studies of the WHIM possible.   Interstellar OVI was first detected in the UV with the {\\em Copernicus} satellite (Jenkins, 1978a, 1978b, and York, 1977), but only recently has data from the {\\em Far Ultraviolet Spectroscopic Explorer} (FUSE) satellite shown the presence of extragalactic OVI intervening absorpion (e.g.  Tripp et al., 2001, Sembach et al., 2000).   In this paper we present the first X-ray detection of highly ionized absorption along the line of sight towards the blazar PKS~2155-304, with the {\\em Chandra} HRCS-LETG ({\\em High Resolution Camera Spectrometer}-{\\em Low Energy Transmission Grating}: Brinkman et al., 2000).  We argue that the physical and dynamical conditions implied by these lines require an extragalactic, low density, origin. Hence these lines are the first detection of the X-ray Forest. ",
        "conclusions": "In this paper we report the first discovery of the ``X-ray Forest'' of absorption lines produced by highly ionized intergalactic medium.  High significance X-ray absorption lines of OVII, OVIII and NeIX are detected along the line of sight to the bright blazar PKS~2155-304, and are associated with a known OVI UV absorber.  We demonstrate that the dynamical properties of the X-ray and the UV absorber are fully consistent with each other, and that a reasonable common and self-consistent physical solution can be found only if photoionization and collisional ionization both contribute to the ionization of the absorbing gas. This requires electron densities of about $6 \\times 10^{-6}$ cm$^{-3}$ for the diffuse X-ray background to be significant. This low denisty requires a linear size, along the line of sight, of the order of $3 [O/H]_{0.3\\times\\odot}^{-1}$ Mpc.  This clearly locates the absorber outside our Galaxy in intergalactic space. We demonstrate that both the dynamical and physical properties of such an absorber are remarkably consistent with those predicted for the low redshift warm phase of the IGM as predicted by hydrodynamical simulations for the formation of structures in the Universe.  Finally we find that solutions with a Ne/O ratio of about 2.5 times solar are favored, suggesting type-II Supernova enrichment, or the presence of dust in the IGM.  \\begin{center} {\\bf Aknowledgements} \\end{center} This work has been partly supported by the NASA grant ADP NAG-5-10814 (FN), the {\\em Chandra} grant DDO-1005X FN and AF), CXC grant NAS8-39073 (FN and AF) and the NASA grant LTSA NAG-5-8913 (SM)."
    },
    "0201/astro-ph0201091_arXiv.txt": {
        "abstract": "Three cometary-shaped objects in the giant \\HII\\ region NGC~3603, originally found and identified as proto-planetary disks (ProPlyDs) by \\cite{brandner2} using HST+VLT in the optical and near-infrared, have been detected with the Australia Telescope Compact Array (\\ATCA\\thanks{The Australia Telescope is funded by the Commonwealth of Australia for operation as a National Facility managed by CSIRO.}) in the radio continuum at 3 and 6~cm. All three {\\plike} objects are clearly resolved with an extent of a few arcseconds. The integrated 6~cm fluxes are up to 1.3 times higher than the 3~cm fluxes with spectral indices averaged over the whole clump between $\\alpha=-0.1$ and --0.5 ($S_\\nu \\propto \\nu^{\\alpha}$), indicating the likely presence of non-thermal emission in at least some of the sources. We present spectral index maps, and show that the sites of negative radio spectral indices are predominantly concentrated in the direction of the tails in at least two of the three {\\plike} nebulae while positive spectral indices are found in the region facing the ionizing star cluster. We propose that thermal bremsstrahlung and non-thermal synchrotron radiation are at work in all three {\\plike} sources. In at least one of the three objects optically thin non-thermal synchrotron emission appears to dominate when averaged over their whole spatial extent, while the spectrum of the second source shows a marginal indication of a non-thermal spectrum. The average spectrum of the third source is in agreement with thermal bremsstrahlung. All measured fluxes are at least one order of magnitude higher than those predicted by \\cite{brandner2}. Upper limits for mass loss rates due to photo-evaporation are calculated to be $\\sim 10^{-5}$~\\Msun\\,year$^{-1}$ and for electron densities to be $\\sim 10^{4}$~cm$^{-3}$. Due to the unexpectedly large radio luminosities of the {\\plike} features and because the radio emission is extended a (proto-)stellar origin of the non-thermal emission from a dust enshrouded star appears unlikely. Instead we propose that magnetized regions within the envelope of the {\\plike} nebulae exist. ",
        "introduction": "The giant \\HII\\ region NGC~3603, located at a distance of about 6~kpc, shows the densest concentration and largest collection of visible massive stars known in our Galaxy. Recent HST images of $\\sim$0\\farcs2 angular resolution \\citep{moffat} have shown that NGC~3603 consists of three Wolf-Rayet (WR) stars and $\\sim$70 O-type stars, with an estimated 40--50 of these stars located in the central $\\sim$30\\arcsec$\\times$30\\arcsec\\ (1 pc $\\times$ 1 pc) region. The three H-rich WR-stars of subtype WNL are also located within $\\sim$1\\arcsec\\ of each other in the central core and are the brightest members of the cluster. They are believed to be massive main sequence stars which, like the bright stars in R136a, drive very strong winds \\citep{dekoter,crowther}. NGC~3603 has also been shown to be a seat of active star formation outside the main cluster region (e.g. \\cite{brandner2}). In particular, three distinct dense cometary nebulae were detected there using HST and VLT: Are these ultra compact \\HII\\  regions (UCHR) or ProPlyDs ?  UCHRs are small ($<$0.1~pc) nebulae, {\\em internally} photoionized by a deeply embedded, massive star. Even though the embedded stars are very luminous, they are invisible at optical wavelengths because of the surrounding dust; instead they show strong far-infrared emission. Their radio continuum emission is often associated with OH and H$_2$O masers. The morphology of UCHRs in high resolution radio continuum images ranges from spherical, cometary, core-halo, shell, to irregular \\citep{wood}.  In contrast, a proto-planetary disk (ProPlyD) is a phenomenon describing a low-mass, young stellar object (YSO) with a circumstellar disk, embedded in a dense (neutral and ionized) envelope that is being {\\em externally} photo-evaporated by the ultraviolet radiation from one or more massive stars. These low-mass stars are not able to ionize their surroundings significantly. The disks in these systems are either observed directly (as in Orion where the disks are seen directly or in silhouette against the bright background nebula \\citep{bally}), or are inferred because in almost all of the ProPlyDs (young) low-mass stars are visible at optical or near-infrared wavelengths (see e.g. \\cite{stecklum}). At low angular resolution ProPlyDs look like UCHRs, but at high resolution, optical and near-infrared images reveal their different nature. So most ProPlyDs would have been classified as UCHRs prior to those observations. In some cases the distinction between the two classes remains difficult. Given the fact that there are about four times more UCHRs than expected \\citep{churchwell} from star formation rates and the potential problems identifying host stars for UCHRs (e.g. some may not have enough FIR flux to account for an internal OB star), it may well be that a large fraction of catalogued UCHRs are mis-identified ProPlyDs. Interferometric radio continuum observations of ProPlyDs are needed to determine their structure, spectral indices and thus the nature of their emission, as well as mass loss rates and extinctions when compared to measured \\Ha\\ fluxes.  ProPlyDs were first identified in the Orion Nebula\\footnote{ There are a variety of different terms used in the literature for  describing the knots of ionised gas in M\\,42 (Orion): e.g. CKs (cometary knots), PIGs (partially ionized globules), EIDERS (external ionised (accretion) disks in the environs of radiation sources), and ProPlyDs (proto-planetary disks); for a summary see \\cite{mccullough}. Another expression being used is EGGs (for evaporating gaseous globules) describing the objects found in M\\,16 (Eagle Nebula). The cometary knots found in the Helix Nebula are compact globules and very different from ProPlyDs as they contain no stars.}, where over 150 of them are known \\citep{odell}. Two ProPlyDs have been identified in more distant nebulae: one in NGC~2024 by \\cite{stapelfeldt} and one, G5.97--1.17, in the Lagoon Nebula by \\cite{stecklum}. The three {\\plike} objects, hereafter referred to P1, P2 and P3, in NGC~3603 \\citep{brandner2} are the biggest, youngest and most massive ones found {\\it{so far}}. They are also the most distant known. These emission nebulae are clearly resolved in the HST/WFPC2 observations, and share the overall morphology of the ProPlyDs in Orion. All three nebulae are rim-brightened and tear-drop shaped with the tails pointing away from the central ionizing cluster.  According to \\cite{brandner2} the brightest object (P1), which has a projected distance of 1.3~pc  from the cluster, has the spectral (excitation) characteristics of a UCHR. Optical spectra reveal the presence of an underlying, heavily reddened continuum source, which is also confirmed by near-infrared VLT/ISAAC observations.  The WFPC2 observations show that only the outermost layer is ionized whereas the interior is neutral. The morphology of P1 is described as a heart-shaped head with a collimated structure in between, which can be understood as the superposition of two individual ProPlyDs. In contrast, P2 and P3, located at projected distances of 2.2~pc and 2.0~pc from the stellar cluster, respectively, show approximately axisymmetric morphologies. No embedded disk or central star has been detected so far in any of these nebulae, preventing a clear identification as proto-planetary disks. The optical point source (see Fig.~2c) close to P3 is probably not physically linked to the nebula \\citep{brandner2}. The {\\plike} structures in NGC~3603 are about two orders of magnitude fainter than typical UCHRs, but have a typical extent of 9000 AU with tails extending to 21000 AU, much larger in size than the ProPlyDs in Orion.  Recent 3.4~cm radio continuum and recombination line measurements of NGC~3603 by \\cite{depree}, which were focussed on abundance measurements and the bright continuum emission from the ionized gas in this region, have an angular resolution of $\\sim$7\\arcsec\\ and a $5\\sigma$ sensitivity of 55 mJy. By obtaining high sensitivity and high angular resolution ($\\sim$1\\arcsec--2\\arcsec) {\\ATCA} observations we primarily aimed to study the radio emission of the winds from many of the early-type stars in the cluster, as well as to detect and resolve the {\\proplyd}s and other gaseous regions in the cluster periphery. This paper will focus on the {\\plike} sources\\footnote{Throughout this paper we use the term ``ProPlyD-like'' for these cometary-shaped objects in NGC~3603, since the identification of these objects as true ProPlyDs is premature, owing to the lack of a clear detection of a central disk or star.} only, which have been detected and are shown to be clearly resolved with the \\ATCA. A subsequent paper containing a detailed study of the whole NGC~3603 region based on our \\ATCA\\ observations will follow. ",
        "conclusions": "The three massive ($\\sim$1--10~\\Msun) {\\plike} nebulae in NGC~3603, which have recently been discovered by \\cite{brandner2}, have been clearly detected and resolved with the {\\it ATCA} at 3 and 6~cm, with one of them likely to be composed of two cometary-shaped objects. Their flux densities are about 10--20 times higher than expected from the \\Ha\\ measurements, and a non-thermal average spectrum can be associated with at least one of the three {\\plike} sources. This is the first time that non-thermal radio emission has been detected from {\\plike} sources. Our spectral index maps show that the emission region is  rather inhomogeneous, with negative spectral indices in the tail and part of the head whereas positive spectral indices, indicating thermal free-free  emission, tend to be detected from a small region facing the star cluster.  We derive upper limits for the mass-loss rates of $10^{-5}$~\\Msun\\,year$^{-1}$ and electron densities of $10^{4}$~cm$^{-3}$. These are in reasonable agreement with estimates from recent HST-images.  We show that synchrotron emission from a magnetized, relativistic, non-thermal particle population may explain the non-thermal spectral regions. Energetic electrons necessary for synchrotron radiation may be produced through shocks or by magnetic reconnection out of the pool of thermal particles, or may be Galactic cosmic rays.  A stellar origin of the observed flux densities appears unlikely considering that the sources are extended. Thus magnetic fields, which have been neglected in ProPlyD models so far, are possibly associated with the ProPlyD envelopes or disks, and appear to be crucial in understanding the physical processes in these {\\plike} objects.  \\cite{melnick} have shown that an inhomogeneous dust distribution in NGC~3603 causes increasing extinction with distance from the star cluster. In particular, the extinction at the location of the {\\proplyd}s is estimated to lie between $A_V \\approx 5-6$ mag. Extinction estimates derived from the radio- to {\\Ha}-luminosity ratio are 1--2 magnitudes higher.  The ultimate discovery of the so-far undetected disks, which are thought to be part of all ProPlyDs, would give important hints about the origin of the extremely high and non-thermal radio fluxes and the nature of the cometary-shaped clumps seen in NGC~3603. This may be accomplished in future mm observations with an upgraded \\ATCA."
    },
    "0201/astro-ph0201328_arXiv.txt": {
        "abstract": "Deep extragalactic surveys with ISOCAM revealed the presence of a large density of faint mid-infrared (MIR) sources. We have computed the 15\\,$\\mu$m integrated galaxy light produced by these galaxies above a sensitivity limit of 50 $\\mu$Jy. It sets a lower limit to the 15\\,$\\mu$m extragalactic background light of (2.4 $\\pm$ 0.5) nW m$^{-2}$ Hz$^{-1}$.  The redshift distribution of the ISOCAM galaxies is inferred from the spectroscopically complete sample of galaxies in the Hubble Deep Field North (HDFN). It peaks around $z\\sim$ 0.8 in agreement with studies in other fields. The rest-frame 15\\,$\\mu$m and bolometric infrared (8-1000\\,$\\mu$m) luminosities of ISOCAM galaxies are computed using the correlations that we establish between the 6.75, 12, 15\\,$\\mu$m and infrared (IR) luminosities of local galaxies. The resulting IR luminosities were double-checked using radio (1.4 GHz) flux densities from the ultra-deep VLA and WSRT surveys of the HDFN on a sample of 24 galaxies as well as on a sample of 109 local galaxies in common between ISOCAM and the NRAO VLA Sky Survey (NVSS). This comparison shows for the first time that MIR and radio luminosities correlate up to $z\\sim$ 1. This result validates the bolometric IR luminosities derived from MIR luminosities unless both the radio-far infrared (FIR) and the MIR-FIR correlations become invalid around $z\\sim$ 1.  The fraction of IR light produced by active nuclei was computed from the cross-correlation with the deepest X-ray surveys from the Chandra and XMM-Newton observatories in the HDFN and Lockman Hole respectively. We find that at most 20\\,$\\%$ of the 15\\,$\\mu$m integrated galaxy light is due to active galactic nuclei (AGNs) unless a large population of AGNs was missed by Chandra and XMM-Newton.  About 75\\,$\\%$ of the ISOCAM galaxies are found to belong to the class of luminous infrared galaxies ($L_{\\rm IR}$ $\\geq$ $10^{11}$ $L_{\\sun}$). They exhibit star formation rates of the order of $\\sim$ 100 $M_{\\sun}$ $yr^{-1}$. The comoving density of infrared light due to these luminous IR galaxies was more than 40 times larger at $z\\sim$ 1 than today.  The contribution of ISOCAM galaxies to the peak of the cosmic infrared background (CIRB) at 140\\,$\\mu$m was computed from the MIR-FIR correlations for star forming galaxies and from the spectral energy distribution of the Seyfert 2, NGC 1068, for AGNs. We find that the galaxies unveiled by ISOCAM surveys are responsible for the bulk of the CIRB, i.e (16 $\\pm$ 5) nW m$^{-2}$ Hz$^{-1}$ as compared to the (25 $\\pm$ 7) nW m$^{-2}$ Hz$^{-1}$ measured with the COBE satellite, with less than 10\\,$\\%$ due to AGNs. Since the CIRB contains most of the light radiated over the history of star formation in the universe, this means that a large fraction of present-day stars must have formed during a dusty starburst event similar to those revealed by ISOCAM. ",
        "introduction": "The extragalactic background light (EBL) is a measurement of the sum of the light produced by all extragalactic sources over cosmic time. When it is integrated over the full spectral range, the so-called cosmic background is a fossil record of the overall activity of all galaxies from their birth until now. It can be considered as the global energetic budget available for any model aiming at simulating the birth and fate of galaxies during the Hubble time. However the physical origin of this light will remain unknown until we have pinpointed the individual sources responsible for it. The goal of the present paper is to demonstrate that an important new result has come from the combination of a series of deep extragalactic surveys performed in the mid-infrared (MIR) at 15\\,$\\mu$m with the ISOCAM camera (Cesarsky et al. 1996a) onboard the Infrared Space Observatory (ISO, Kessler et al. 1996): we suggest here that the galaxies detected in these surveys, which median redshift of $z\\sim 0.8$ was measured from a sub-sample of ISOCAM galaxies, contribute dominantly to the cosmic infrared background (CIRB), i.e. the EBL integrated over all wavelengths within $\\lambda$= 5 to 1000\\,$\\mu$m.  The CIRB was recently detected and measured thanks to the cosmic background explorer (COBE) instruments FIRAS (Far Infrared Absolute Spectrometer) and DIRBE (Diffuse Infrared Background Experiment) (Puget et al. 1996, Fixsen et al. 1998, Lagache et al.  1999, 2000, Hauser et al. 1998, Dwek et al. 1998, Finkbeiner et al. 2000) from 100\\,$\\mu$m to 1 mm. It peaks around $\\lambda_{max}\\simeq$ 140\\,$\\mu$m and was found to represent at least half and maybe two thirds of the overall cosmic background (see Gispert, Lagache \\& Puget 2000). Hence the CIRB reflects the bulk of the star formation that took place over the history of the universe. By resolving it into individual galaxies, we would therefore pinpoint the times and places where most stars seen in the local universe were formed.  Two physical processes were considered for its origin: nucleosynthesis, i.e. stellar radiation in star forming galaxies, and accretion around a black hole, i.e. active galactic nuclei. In both cases, the light is not directly coming from its physical source but is reprocessed by dust, i.e. absorbed and re-radiated thermally by the ``warm'' dust. Both processes are probably related (see Genzel et al. 1998), but energetic considerations, based on the presence of massive black holes and on the amount of heavy elements in local galaxies, suggest that star formation should by far dominate in the CIRB over AGN activity (Madau \\& Pozzetti 2000, Franceschini et al. 2001). However, until the individual galaxies responsible for the CIRB are found and studied in detail, this result will remain theoretical.  The spectral energy distribution (SED) in the IR of local galaxies peaks above $\\sim$60\\,$\\mu$m and typically around 80 $\\pm$ 20\\,$\\mu$m (see Sanders \\& Mirabel 1996). As a result, the distant galaxies responsible for the peak of the CIRB detected by COBE around $\\lambda_{\\rm max}$ $\\sim$ 140\\,$\\mu$m should be located below $z\\sim$ 1.3 and present a redshift distribution peaked around $z\\sim$ 0.8, if their SEDs do not strongly differ from those of local galaxies. As we will see this is also the redshift range of the galaxies detected at 15\\,$\\mu$m with ISOCAM.  The ISOCAM extragalactic surveys were performed with two filters, LW2 (5-8.5\\,$\\mu$m) and LW3 (12-18\\,$\\mu$m), centered at 6.75 and 15\\,$\\mu$m respectively. The 6.75\\,$\\mu$m sample of sources is strongly contaminated by galactic stars, whereas stars are rather easily distinguished from galaxies at 15\\,$\\mu$m using optical-MIR colour-colour plots. As a consequence, we are only concerned here by the 15\\,$\\mu$m galaxies. Moreover, the observed 6.75\\,$\\mu$m light is no more produced by dust emission for galaxies more distant than $z\\sim$ 0.4 because of k-correction (redshifted stellar light dominates the 6.75\\,$\\mu$m band above this redshift), whereas the observed 15\\,$\\mu$m light is mostly due to dust emission for galaxies up to $z\\sim$ 2.  About 1000 galaxies detected in the 15\\,$\\mu$m surveys were used to produce number counts (i.e. surface density of galaxies as a function of flux density; see Elbaz et al. 1999). The steep slope of the 15\\,$\\mu$m counts below $\\sim\\,1$ mJy indicates the presence of an excess of faint sources by one order of magnitude in comparison with predictions assuming no evolution of the 15\\,$\\mu$m luminosity function with redshift. The presence of broad emission features in the MIR spectrum of galaxies alone cannot explain the shape of the number counts and a strong evolution of either the whole luminosity function (Xu 2000, Chary \\& Elbaz 2001) or preferentially of a sub-population of starburst galaxies evolving both in luminosity and density (Franceschini et al. 2001, Chary \\& Elbaz 2001, Xu et al. 2001) is required in order to fit the ISOCAM 15\\,$\\mu$m counts. In the present paper, we suggest that these ISOCAM galaxies are in fact dusty starbursts responsible for the bulk of the CIRB.  In Sect.~\\ref{resolve}, we compare the sensitivity of different extragalactic surveys in several wavelength ranges to detect the galaxies responsible for the CIRB. It is suggested that MIR is presently the most efficient technique to detect dusty starbursts up to $z\\sim$ 1.3.  In Sect.~\\ref{ebl15}, we calculate the 15\\,$\\mu$m integrated galaxy light (IGL) due to ISOCAM galaxies. The 15\\,$\\mu$m IGL is the sum of the 15\\,$\\mu$m fluxes from individual galaxies, down to a given sensitivity limit, per unit area. It represents a lower limit to the 15\\,$\\mu$m EBL, which remains unknown. Once the redshift distribution and SED of these galaxies is determined, it becomes possible to estimate their contribution to the CIRB.  In Sect.~\\ref{MIRtracer}, we demonstrate that MIR luminosities at 6.75, 12 and 15\\,$\\mu$m are strongly correlated with the bolometric IR luminosity (from 8 to 1000\\,$\\mu$m) for local galaxies. The correlations presented in Chary \\& Elbaz (2001) are confirmed here with a larger sample of galaxies.  Before spectroscopic redshifts are obtained for the full sample of ISOCAM galaxies used to produce these number counts, the redshift distribution of these galaxies can be inferred from a few sub-samples: HDFN (Aussel et al. 1999, 2001), CFRS-14 (Flores et al. 1999, 2002), CFRS-03 (Flores et al. 2002). The ultra-deep ISOCAM survey of the HDFN samples a flux density range where most of the evolution observed in the number counts takes place and where the bulk of the 15\\,$\\mu$m IGL is produced. This field is complete in spectroscopic redshifts, so it is used to estimate the bolometric IR luminosities and star formation rates of the ISOCAM galaxies in Sect.~\\ref{bolIR}.  This result relies on two assumptions:\\\\ - that the main source for the MIR light in ISOCAM galaxies is star formation and not accretion around a black hole.\\\\ - that the correlations found in the local universe between the MIR and bolometric IR luminosity of galaxies remain valid up to $z\\sim$ 1.  The first assumption is discussed and justified in Sect.~\\ref{agnfrac}, where soft and hard X-ray data from the Chandra and XMM-Newton X-ray observatories are combined with ISOCAM data on galaxies in the HDFN and Lockman Hole regions respectively.  The issue of the robustness of the MIR-FIR correlations in the distant universe is addressed in Sect.~\\ref{radio}, where IR luminosities are also computed from radio (1.4 GHz) flux densities for a sub-sample of 24 distant and 109 local ISOCAM galaxies.  In Sect.~\\ref{ldens}, we compute the cosmic density of IR light due to luminous IR galaxies ($L_{\\rm IR}\\geq 10^{11}~L_{\\sun}$) at ${\\bar z}\\sim$ 1. In Sect.~\\ref{cirb}, we evaluate the contribution of the ISOCAM galaxies to the CIRB, more precisely to its peak emission around $\\lambda_{max}\\sim$ 140\\,$\\mu$m. Finally, the nature of ISOCAM galaxies is discussed in the conclusions (Sect.~\\ref{conclusions}).  In the following, we will use the terms ULIG for galaxies with an IR luminosity $L_{\\rm IR}=~L[8-1000\\,\\mu{\\rm m}]\\geq 10^{12}~L_{\\sun}$, LIG, when $10^{11}\\leq (L_{\\rm IR}/L_{\\sun})< 10^{12}$ and luminous IR galaxies for both ($L_{\\rm IR}\\geq~10^{11}~L_{\\sun}$).  Throughout this paper, we will assume H$_o$= 75 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\rm matter}$= 0.3 and $\\Omega_{\\Lambda}= 0.7$. ",
        "conclusions": "\\label{conclusions} The cosmic IR background is a fossil record of the light radiated by galaxies since their formation. It peaks around $\\lambda_{\\rm max}\\simeq$ 140\\,$\\mu$m whereas the spectral energy distribution of galaxies peaks above 60\\,$\\mu$m. This suggests that the bulk of the cosmic IR background is due to galaxies located below $z\\sim$ 1.3. We have shown that the best technique currently available to unveil dusty galaxies up to $z\\sim$ 1.3 is provided by the ISOCAM MIR extragalactic surveys.  We have computed the contribution of ISOCAM galaxies to the 15\\,$\\mu$m background, the 15\\,$\\mu$m integrated galaxy light, and found a value of $IGL_{15}\\simeq$ (2.4$\\pm$0.5) nW m$^{-2}$ sr$^{-1}$. This is about ten times below the cosmic background measured by COBE at $\\lambda_{\\rm max}\\simeq$ 140\\,$\\mu$m: $EBL_{140}\\simeq$ (25$\\pm$7) nW m$^{-2}$ sr$^{-1}$.  We have demonstrated that the MIR luminosities at 6.75, 12 and 15\\,$\\mu$m were correlated with each other and with the bolometric IR luminosity for local galaxies. This suggests that the contribution of ISOCAM galaxies to the CIRB can be computed from $IGL_{15}$, unless distant galaxies SEDs strongly differ from local ones.  The redshift distribution of ISOCAM galaxies was measured from the spectroscopically complete sample of galaxies in the region of the HDFN. This redshift distribution is consistent with the twice larger sample of ISOCAM galaxies detected in the CFRS fields CFRS-14 and CFRS-03 (Flores et al. 1999, 2002). At the median redshift of ${\\bar z}\\simeq$ 0.8, the observed 15 and 140\\,$\\mu$m wavelengths correspond to about 7\\,$\\mu$m (ISOCAM-LW2 filter) and 80\\,$\\mu$m (IRAS bands) in the rest-frame of the galaxies. Luminosities at both wavelengths are correlated (see Fig.~\\ref{FIG:correl}d). If the correlations between MIR and FIR luminosities remain valid up to $z\\sim$ 1, then they can be used to compute $IGL_{140}$.  We have checked with a sample of galaxies detected both in the MIR with ISOCAM and in the radio with the VLA and WSRT, that the MIR-FIR and radio-FIR correlations are consistent up to $z\\sim$ 1.  This comparison independently validates our estimate of the bolometric IR luminosity of the ISOCAM galaxies, although it is not clear whether the radio-FIR correlation works also up to $z\\sim$ 1.  The fraction of active nuclei responsible for the 15\\,$\\mu$m luminosity of ISOCAM galaxies was estimated from the deepest soft and hard X-ray surveys available at present by the XMM-Newton and Chandra X-ray observatories in the Lockman Hole and HDFN respectively (Fadda et al. 2001). It was found that about (12$\\pm$5)\\,$\\%$ of the ISOCAM galaxies are powered by an AGN and that the AGN contribution to $IGL_{15}$ was about (17$\\pm$6)\\,$\\%$. The AGN contribution to $IGL_{140}$ was found to be as low as $\\sim$ 4\\,$\\%$ assuming that they all share the SED of the local Seyfert 2, NGC 1068. This is a conservative choice since NGC 1068 presents the flattest IR SED that we know. However, we note that the cosmic X-ray background (CXB) peaks around 30 keV (see Fig.1 of Wilman, Fabian \\& Nulsen 2000), while both XMM-Newton and Chandra were limited to energies below 10 keV. It is therefore possible that a population of hard X-ray AGNs was missed by these surveys. But as claimed by the authors of these deep X-ray surveys, the bulk of the CXB had been resolved into individual galaxies in the Lockman Hole and HDFN images. Moreover, using estimates of the present comoving density of black holes, Madau \\& Pozzetti (2000) calculated that less than 20\\,$\\%$ of the CIRB could be due to dusty AGNs.  For the remaining star forming galaxies, we used a library of template SEDs, reproducing the MIR-FIR correlations, to compute their IR luminosity and contribution to $IGL_{140}$. We find that LIGs and ULIGs produce about 60\\,$\\%$ of $IGL_{15}$. The comoving density of IR luminosity produced by these luminous IR galaxies was about (70$\\pm$35) times larger at $z\\sim$ 1 than today, while in the same redshift interval the $B$-band or UV luminosity densities only decreased by a factor $\\sim$ 3. Since the IR luminosity measures the dusty star formation rate of a galaxy, this also implies that the comoving density of star formation, due to luminous IR galaxies, decreased by a similar factor in this redshift range, i.e. much more than expected by studies at other wavelengths affected by dust extinction. This indicates that a large fraction of present day stars were formed during a dusty starburst event.  We estimate a contribution of ISOCAM galaxies to the peak of the CIRB at $\\lambda_{max}\\simeq$ 140\\,$\\mu$m of (16$\\pm$5) nW m$^{-2}$ sr$^{-1}$ as compared to the measured value of (25$\\pm$7) nW m$^{-2}$ sr$^{-1}$ from COBE. This study therefore suggests that the ISOCAM galaxies are responsible for the bulk of the CIRB.  We have started a systematic spectroscopic follow-up of these galaxies with the aim of studying their physical properties and the origin of their large SFR. Franceschini et al. (2001) estimated their baryonic masses to be of the order of $\\left< M \\right> \\sim 10^{11}~M_{\\sun}$ by fitting their optical and near-IR luminosities with template SEDs (from Silva et al. 1998) and assuming a Salpeter initial mass function (from 0.15 to 100\\,$M_{\\sun}$). Their colors are similar to field galaxies of similar magnitudes (Cohen 2001), hence they could not have been selected on the basis of their optical colors. The technique which consists in using the spectral slope in the UV domain to correct luminosities from extinction (Meurer, Heckman \\& Calzetti 1999) fails to differentiate the luminous dusty galaxies detected with ISOCAM from other field galaxies in the HDFN+FF (Cohen 2001). This was to be expected since this technique only works for galaxies with $L_{\\rm IR}\\lesssim~4\\times10^{11}~L_{\\sun}$ (Meurer et al. 2000) while most ISOCAM galaxies are more luminous than this threshold. A property of the ISOCAM galaxies that may give a hint on their origin is their association with small groups of galaxies. A preliminary study of 22 ISOCAM-HDFN galaxies by Cohen et al. (2000) found that nearly all ISOCAM galaxies belonged to small groups, while the fraction of field galaxies with similar optical magnitudes belonging to such groups was 68\\,$\\%$. The study of the full sample of ISOCAM-HDFN galaxies by Aussel et al. (2001) shows that all of them belong to physical groups, hence suggests that dynamical effects such as merging and tidal interactions are responsible for their large SFR.  The Space IR Telescope Facility (SIRTF) will soon provide a powerful insight on the FIR emission of distant dusty starbursts and its 24\\,$\\mu$m band, less affected by confusion, should prolongate their detection up to $z\\sim2.5$. Extending the redshift range surveyed by ISOCAM to such redshifts is particularly important to measure the direct IR emission of the distant population of Lyman break galaxies, whose dust extinction remains highly uncertain (Steidel et al. 1999). However, the direct measurement of the MIR and FIR emission of distant galaxies will only become possible with the launch of the Next Generation Space Telescope (NGST) and Herschel (FIRST) satellite scheduled for 2009 and 2007."
    },
    "0201/astro-ph0201434_arXiv.txt": {
        "abstract": "{ Equilibrium configurations of cold neutron stars near the minimum mass are studied, using the recent equation of state SLy, which describes in a unified, physically consistent manner, both the solid crust and the liquid core of neutron stars. Results are compared with those obtained using an older FPS equation of state of cold catalyzed matter. The value of $M_{\\rm min}\\simeq 0.09~{\\rm M}_\\odot$ depends very weakly on the equation of state of cold catalyzed matter: it is $0.094~{\\rm M}_\\odot$ for the SLy  model, and $0.088~{\\rm M}_\\odot$ for the FPS one. Central density at $M_{\\rm min}$ is significantly lower than the normal nuclear density: for the SLy equation of state we get central density $1.7~10^{14}~{\\rm g~cm^{-3}}$, to be compared with $2.3~10^{14}~{\\rm g~cm^{-3}}$ obtained for the FPS one. Even at $M_{\\rm min}$, neutron stars have a small liquid core of radius of about 4 km, containing some   2-3\\% of  the stellar mass. Neutron stars with $0.09~{\\rm M}_\\odot<M<0.17~{\\rm M}_\\odot$ are bound with respect to dispersed configuration of the hydrogen gas, but are unbound with respect to dispersed $^{56}{\\rm Fe}$. The effect of uniform rotation on the minimum-mass configuration of cold neutron stars is studied. Rotation increases the value of $M_{\\rm min}$; at rotation period of 10 ms  the minimum mass of neutron stars increases to $0.13~{\\rm M_\\odot}$, and corresponds to the mass-shedding (Keplerian) configuration. In the case of the shortest observed rotation period of radio pulsars  1.56 ms, minimum mass of uniformly rotating cold neutron stars corresponds to the mass-shedding limit, and is found at $0.61~{\\rm M}_\\odot$ for the SLy EOS and $0.54~{\\rm M}_\\odot$ for the FPS EOS. ",
        "introduction": "Mass of neutron stars is bound on the lower-density side by the minimum mass, $M_{\\rm min}$. Equilibrium neutron star configurations with $M<M_{\\rm min}$ do not exist. Some astrophysical scenarios, in which a neutron star looses matter, and reaches the critical value $M_{\\rm min}$, were proposed (Blinnikov et al. 1984, Colpi et al. 1989, Colpi et al. 1991,  Sumiyoshi et al. 1998). Basically, one considers a neutron star in a close binary system with a more compact, and more massive object (more massive neutron star, stellar mass black hole). At a sufficiently small separation between a neutron star and its companion, less massive, and therefore larger, neutron star starts to lose mass  due to gravitational pull of its companion, the lost matter being accreted by the companion compact object. This process is self-accelerating, because decrease of mass leads to increase of neutron star size, making it even more susceptible to the mass loss. After crossing $M_{\\rm min}$, no equilibrium configuration can be reached. Numerical simulations suggest, that after crossing the $M_{\\rm min}$ value, a neutron star will undergo an explosion (Blinnikov  et al. 1984, Colpi et al. 1989, 1991, Sumiyoshi et al. 1998).  Of course, rotation and thermal effects can affect (actually - increase) the actual value of the minimum neutron star mass. As we will see, rotation will also change the nature of the minimum-mass  configuration. In the present paper, we will restrict ourselves to neutron stars built of cold catalyzed matter. For newly born  proto-neutron stars, both thermal and neutrino-trapping effects are large, and are expected to increse minimum mass to $0.9-1.1~{\\rm M}_\\odot$ (Goussard et al. 1998, Strobel et al. 1999). However, $M_{\\rm min}$ for static, cold neutron stars still remains the absolute lower bound on the mass of neutron star.  The calculation of $M_{\\rm min}$ for neutron stars has a long history. The first correct estimate of $M_{\\rm min}$ was obtained by Oppenheimer \\& Serber (1938). They showed, that the previous estimate obtained by Landau (1938) was much too small, and moreover based on incorrect arguments. Oppenheimer \\& Serber (1938) obtained an estimate $0.17~{\\rm M}_\\odot$ without taking into account nuclear interactions. Then they argued, that inclusion of the contribution from the nucleon-nucleon interactions leads to a decrease of $M_{\\rm min}$ to $0.03-0.10~ {\\rm M}_\\odot$, depending on the assumptions they made about the (very poorly known  at that time) nucleon-nucleon interaction. It should be stressed, that their criterion of stability was based  on energy arguments: a neutron star had to be stable with respect to transformation into medium mass nuclei (they took calcium as an example). Calculation of Harrison et al. (1964) was based on the Harrison-Wheeler equation of state (Harrison et al. 1964), which described both the crust and the liquid core of neutron stars. From the present-day point of view, their equation of state above neutron drip was unrealistic. They obtained $M_{\\rm min}=0.18~{\\rm M}_\\odot$, at a central density of only $3~10^{13}~{\\rm g~cm^{-3}}$. Using two versions of the Levinger \\& Simmons (1961) baryon-baryon potential, Tsuruta \\& Cameron (1966) obtained $M_{\\rm min}=0.11-0.13~{\\rm M}_\\odot$. Cohen \\& Cameron (1971), using the equation of state (EOS) based on the up-dated version of the Levinger-Simmons nucleon-nucleon potential (Langer et al. 1969), obtained $M_{\\rm min}=0.065~{\\rm M}_\\odot$. Finally, in their classical paper on the EOS and structure of neutron stars, Baym et al. (1971b, hereafter BPS), who used their own (BPS) EOS of the outer crust, and that of Baym et al. (1971a, hereafter BBP) for the inner crust and for the outer layers of the liquid core, obtained $M_{\\rm min}=0.0925~{\\rm M}_\\odot$.   The properties of the equilibrium configurations  near $M_{\\rm min}$ are sensitive to the EOS of neutron star matter at subnuclear densities, i.e. for mass density significantly below normal (saturation) density of nuclear matter, $\\rho_0=2.7~10^{14}~{\\rm g~cm^{-3}}$ (baryon density significantly below $n_0=0.16~{\\rm fm^{-3}}$). As we will show, a particularly important role is played by the EOS near the inner crust - liquid core boundary. Therefore, the EOS should give physically correct description of this transition region; a brutal {\\it ad hoc} matching of the inner crust and liquid core are not sound for this purpose. The microscopic model, underlying the EOS, should be the same on both sides of the crust-core boundary, and should therefore be based on the same effective nuclear hamiltonian.  In the present paper, we calculate $M_{\\rm min}$ and study static configurations near minimum mass. We also study the effect of rotation on these configurations. We use two EOS, based on different effective nuclear hamiltonians. These effective hamiltonians are:  a recent SLy model (Chabanat et al. 1998, for applications see Douchin \\& Haensel 2000, Douchin et al. 2000, EOS in   Douchin \\& Haensel 2001), and a somewhat older FPS model (Ravenhall \\& Pandharipande 1989, for applications see Lorenz et al. 1993).  The plan of the paper is as follows. Equation of state near crust-liquid core interface is discussed in Sect.2. In Sect. 3, we study equilibrium configurations near $M_{\\rm min}$, and their dependence on the assumed EOS model. The problem of binding energy and stability of low-mass neutron stars is studied in some detail in Sect.4. Possible effects of elastic strain within the equilibrium configurations near $M_{\\rm min}$ are estimated in Sect. 5. Effects of rotation are studied in Sect. 6. Finally, Sect. 7 contains discussion of our results and  conclusions. ",
        "conclusions": "In the present paper we determined the properties of the minimum mass configuration of neutron stars, built of cold catalyzed matter. We used two recent EOS of cold catalyzed matter, which describe in a unified way both the crust and the liquid core, as well as their interface. For static neutron stars, we get a minimum mass of $0.09~ {\\rm  M}_\\odot$,  coinciding with the value of $M_{\\rm min}$ obtained three decades ago in the BPS paper. The value of $M_{\\rm min}$ for neutron stars  built of cold catalyzed matter is quite independent of the specific EOS used.  However, the structure of the minumum mass configuration is sensitive to the EOS at $\\rho\\la \\rho_0$. For the  FPS and SLy EOS we get  $R(M_{\\rm min})\\simeq 220$~km and 270 km, respectively. Our minimum mass configurations contain a tiny liquid core with $2-3\\%$ of the stellar mass. The structure of the $M_{\\rm min}$ configurations for the EOS considered in the present paper are substantially different from that obtained in the classical BPS paper. While BPS got a similar value of $M_{\\rm min}$, the radius of their minimum mass configuration was only 160 km, and  it was completely solid. These differences can be easily understood in terms  of the differences in the EOS at subnuclear density. Namely, for $\\rho\\la \\rho_0$, the BPS EOS is softer, and  its crust-core transition takes place at significantly higher density than for the SLy and FPS ones.  Within mass interval  $M_{\\rm min}=0.09<M<0.17~{\\rm M}_\\odot$, neutron stars are unbound with respect to dispersed $^{56}{\\rm Fe}$. In the similar mass range, neutron stars are less bound than  carbon-oxygen white dwarfs of the same mass. However, all neutron stars are bound with respect to dispersed hydrogen.   Effects of rotation have been shown to be important for $M_{\\rm min}$, and for the low-mass neutron stars in general. For neutron stars rotating rigidly at $P=1.56~$ms, minimum mass is determined by the mass-shedding limit: it is $0.5~{\\rm M}_\\odot$ and $0.6~{\\rm M}_\\odot$, for the FPS and SLy EOS, respectively. The minimum mass configuration at $P=1.56~$ms is quite  oblate, due to strong deformation of the crust by the centrifugal force."
    },
    "0201/gr-qc0201080_arXiv.txt": {
        "abstract": "In this paper we consider the possibility of measuring the corrections induced by the square of the parameter $a_g$ of the Kerr metric to the general relativistic deflection of electromagnetic waves and time delay in an Earth based experiment. It turns out that at laboratory scale the rotational effects exceed definitely the gravitoelectric ones which are totally negligible. By using a small rapidly rotating sphere as gravitating source on the Earth the deflection of a grazing light ray amounts to $10^{-13}$ rad and the time delay is proportional to $10^{-23}$ s. These figures are determined by the upper limit in the attainable values of $a_g$ due to the need of preventing the body from exploding under the action of the centrifugal forces. Possible criticisms to the use of the Kerr metric at $a_g^2$ level are discussed. ",
        "introduction": "The effect of the proper angular momentum $J$ of a body of mass $M$, assumed as source of the gravitational field, is accounted for, in General Relativity, by the parameter $a_g=\\rp{J}{Mc}$, where $c$ is the speed of light in vacuum, entering, e.g., the Kerr metric. It has the dimensions of a length and is of order $\\mathcal{O}(c^{-1})$. The other characteristic length is the Schwarzschild radius $R_s=\\rp{2GM}{c^2}$, where $G$ is the Newtonian gravitational constant, which enters both Schwarzschild and Kerr metrics.  Traditionally, in almost all the applications of the Kerr metric to the motion of test particles or electromagnetic waves in the space--time of a central spinning mass the square of $a_g$, which is of order $\\mathcal{O}(c^{-2})$ and enters the diagonal components of the metric tensor, is neglected. However, in a recent stimulating letter in ref. \\ct{ref:tart02} it is suggested that, instead, it would be better to account for it in view of possible experimental setups on the Earth. Indeed, at laboratory scale, while the effects due to $R_s$, of order $\\mathcal{O}(c^{-2})$, are certainly negligible because of the smallness of the product $GM$, those due to $a_g^2$, of order $\\mathcal{O}(c^{-2})$ as well, could become interesting because, for certain symmetric geometries of the source body, it depends only on the fourth power of its radius $R$ and the square of its angular velocity $\\Omega$.  In this paper we want to explore this scenario by working out the contributions of $a_g^2$ to two classical tests of General Relativity: the deflection of the electromagnetic waves and the time delay \\ct{ref:ciuwhe95}. We will consider the motion of photons in the equatorial plane of a central rotating body at rest. In a number of papers these general relativistic effects have been worked out for various systems of gravitating bodies endowed with mass--monopole, spin--dipole and time--dependent mass and spin multipoles for different states of motions both of the source of electromagnetic waves and the observer \\ct{ref:kop97, ref:blanetal01, ref:kopmash, ref:ser02, ref:sercard02, ref:ciufricc02} by means of a variety of general mathematical approaches, but the influence of $a_g^2$ has never been considered.  A possible objection to the present calculations could be that they are based on the use of the Kerr metric in representing the gravitational field of an extended object at the $a_g^2$ order. Indeed, it should be reminded that, up to now, nobody has been successful in extending the Kerr metric from the empty space to the interior of a body. In fact, the field of a real rotating body would be endowed of various multipolar mass moments, etc. and if we do not know what is the matter source, as in the case of the Kerr metric, we do not know how to cope with them\\footnote{ However, in the present case we could argue that, whatever the description of the non--sphericity of the central mass could be, it should not affect in any relevant way the obtained results, as confirmed a posteriori by the presented calculations.}. Then, one could wonder if, at the level of approximation considered, the Kerr metric and the derived consequences have physical meaning. However, notice that in ref. \\ct{ref:tart98} the author use the full Kerr metric in deriving the general relativistic corrections of the Sagnac effect. In ref. \\ct{ref:tartrug02a} the Kerr metric is used in order to sketch a possible Earth based laboratory experiment aimed to testing gravitomagnetism at order $a_g$. Moreover, recently some efforts are or will be directed (M. McCallum, private communication, 2002) towards possible extensions of the Kerr metric to various interior matter sources. Many efforts have been dedicated to the construction of sources which could represent some plausible models of, e.g., stars. In ref. \\ct{ref:harttorn68} the authors have tried to connect the external Kerr metric to the multipolar structure of various types of stars assumed to be rigidly and slowly rotating. Terms of greater than the second order in the angular velocity were neglected. In ref. \\ct{ref:ali01} a physically reasonable fluid source for the Kerr metric has been obtained. For previous attempts to construct sources for the Kerr metric see refs. \\ct{ref:haggmar81,ref:hagg90}. In ref.\\ct{ref:wilt03} a general class of solutions of Einstein's equations for a slowly rotating fluid source, with supporting internal pressure, is matched to the Kerr metric up to and including first order terms in angular speed parameter. So, in this context, our calculations could be useful in order to give an idea of what could happen in this case and, more generally, if experiments testing different scenarios at order $\\mathcal{O}(J^2)$ are realistically conceivable.  A possible subject for further analysis could be an investigation of the correspondence with the Post--Newtonian expansion at the $J^2$ level. At the first order in $J$ the Post--Newtonian expansion accounts for the angular momentum of the source in the off--diagonal components $g_{0i}\\ i=1,2,3$ via a $c^{-3}$ term\\footnote{This may explain the fact that Kerr metric is commonly accepted in describing the Lense--Thirring effect and the gravitomagnetic clock effects which are linear in $J$. }. So, at a first glance, it could be guessed that the Post--Newtonian expansion should account for $J^2$ with a $c^{-6}$ term which, instead, is absent in the Kerr metric. Indeed, as we will see, our effects depends only on $a_g^2$.  Recently, in ref. \\ct{ref:tart03} an approximated solution of the Einstein field equations for a rotating, weakly gravitating body has been found. It seems promising because both external and internal metric tensors have been consistently found, together an appropriate source tensor. Moreover, the mass and the angular momentum per unit mass are assumed to be such that the mass effects are negligible with respect to the rotation effects. As a consequence, only quadratic terms in the body's angular velocity has been retained. One main concern is that it is not clear if there is a real mass distribution able to generate the found source tensor.  As can be seen, the subject seems to be rather open and, in the author's opinion, worth of investigation, at least in order to get some estimates of the orders of magnitude involved.  The paper is organized as follows: in section 2 we will derive the geodesic motion of a test particle of mass $m$ in the equatorial plane of a rotating body. In section 3 we will specialize it to the photons and will work out the deflection of electromagnetic waves due to $a_g^2$. In section 4 its effect on the time delay is examined. Section 5 is devoted to the discussion of the obtained results and to possible applications to Earth laboratory experimental scenarios. ",
        "conclusions": "Here we investigate the two relativistic effects worked out and examine the possibility of measure them in terrestrial experiments\\footnote{It is important to notice that in an astrometric scenario, e.g., at the limb of the Sun or of a compact star, the gravitoelectric term of order $\\mathcal{O}(c^{-2})$ and the gravitomagnetic term of order $\\mathcal{O}(c^{-3})$ are far larger than the terms due to $a_g^2$ and the approximations used until now are no longer valid. See, e.g., the refs. \\ct{ref:ricmat82a, ref:ricmat82b, ref:ricmat83, ref:bertgiamp92}.}. \\subsection{The deflection of light} \\Rfr{angolo} shows many interesting features \\begin{itemize} \\item  It is of order $\\mathcal{O}(c^{-2})$, as the well known gravitoelectric deflection angle \\ct{ref:ciuwhe95} \\eqi\\delta_{\\rm GE}=2\\rp{R_s}{b}=4\\rp{GM}{c^2 b}\\eqf \\item It is independent of the frequency of the electromagnetic radiation, as $\\delta_{\\rm GE}$ \\item It is opposite in sign with respect to $\\delta_{\\rm GE}$ acting as a diverging effect \\item  It does not contain neither the Newtonian gravitational constant $G$ nor the mass $M$ of the central body, a feature that may be of great help in an Earth laboratory experiment \\item It depends on $b^{-2}$, contrary to $\\delta_{\\rm GE}$ which, instead, depends on $b^{-1}$. \\item  It is insensitive to the sense of rotation of the central body: it is a pity since such a feature could have been useful in order to generate some particular signature. \\end{itemize}  Let us investigate what could happen in a possible Earth laboratory experiment. E.g., if we assume as central body a small sphere of 2.5 cm radius, mass of 111 g and spinning at\\footnote{ Indeed, according to ref. \\ct{ref:tart02b}, the maximum peripheral linear speed attainable by a spinning body is $v_{max}\\sim 10^3$ m s$^{-1}$. } $4\\times 10^4$ rad s$^{-1}$ we would have \\bar R_s^{\\rm sphere} &=& 1.647\\times 10^{-28}\\ \\textrm{m},\\\\ a_g^{\\rm sphere} &=& 3.61\\times 10^{-8}\\ \\textrm{m}. \\ear For a grazing light ray, which means that \\eqi\\delta_{a_g^2}=-\\rp{1}{50}\\left(\\rp{\\Omega R}{c}\\right)^2=-\\rp{1}{50}\\left(\\rp{v_{\\rm per}}{c}\\right)^2,\\eqf the deflection would amount to \\bar \\delta_{\\rm GE}^{\\rm sphere} & = &1.31\\times 10^{-26}\\ \\textrm{rad}=2.717\\times 10^{-21}\\ \\textrm{asec},\\\\ \\delta_{a_g^2}^{\\rm sphere} & = & -2.606\\times 10^{-13}\\ \\textrm{rad}=-5.37\\times 10^{-2} \\ \\mu\\textrm{asec}.\\ear In this case the deflection due to $a_g^2$ is 13 orders of magnitude larger than the gravitoelectric deflection. \\subsection{The time delay} \\Rfr{tre} exhibits the following characteristics \\begin{itemize} \\item  It is of order $\\mathcal{O}(c^{-3})$, as $\\Delta t_{\\rm GE}$ \\item It is independent of the frequency of the electromagnetic radiation, as $\\Delta t_{\\rm GE}$ \\item  It does not contain neither the Newtonian gravitational constant $G$ nor the mass $M$ of the central body, contrary to $\\Delta t_{\\rm GE}$ \\item Its amplitude depends on $b^{-1}$, contrary to $\\Delta t_{\\rm GE}$ which depends only on the characteristic length $R_s$. \\item  It is insensitive to the sense of rotation of the central body \\end{itemize}  At laboratory scale, for the same sphere as before, we have \\bar \\Delta t ^{\\rm sphere}_{\\rm GE} & \\propto & 5.49\\times 10^{-37}\\ {\\rm s },\\\\ \\lb{tdgel} \\Delta t ^{\\rm sphere}_{a_g^2} & \\propto & -8.7\\times 10^{-23}\\ {\\rm s }.\\lb{dtgml}\\ear Also for the time delay the correction due to $a_g^2$ is 14 orders of magnitude larger than the gravitoelectric one. However, the detection of a time difference of the order of $10^{-23}$ s is presently out of discussion: suffice it to say that $10^{-23}$ s is the typical lifetime of a strongly decaying particle (e.g. $\\rho$ or $N^*$). \\subsection{Conclusion} In this paper we have calculated the effect of the square of the parameter $a_g$ entering the Kerr metric on the general relativistic deflection of electromagnetic waves and the time delay in the gravitational field of a central rotating source at rest.  Then, we have analysed those two effects in a possible experimental scenario in an Earth based laboratory with a small spinning sphere.  Due to their extreme smallness, we have neglected the terms involving $GM$. The maximum attainable linear velocity $v_{\\rm max}\\sim 10^3 $ m s$^{-1}$, due to the effects of the centrifugal forces, puts severe upper constraints to the magnitude of $a_g$ in a laboratory experiment. The derived effects are to small to be detected.  The calculations performed here loose their validity in an astrophysical context because, in this case, neglecting the other terms proportional to $GM$ with respect to $a_g^2$ is no more allowed."
    },
    "0201/gr-qc0201049_arXiv.txt": {
        "abstract": "In  this paper, we   analyze the   conditions for  convergence  toward General   Relativity of scalar-tensor  gravity  theories defined by an arbitrary coupling function $\\alpha$ (in the  Einstein frame). We show that, in  general, the  evolution  of the scalar field  $(\\varphi)$ is governed by two opposite  mechanisms:  an attraction  mechanism  which tends to  drive scalar-tensor models toward   Einstein's theory, and a repulsion mechanism  which  has the  contrary effect.  The  attraction mechanism dominates the recent  epochs  of the universe evolution  if, and  only if,  the  scalar field   and its derivative  satisfy certain boundary conditions. Since   these conditions for convergence   toward general relativity depend on  the particular scalar-tensor theory used to describe the universe evolution,  the nucleosynthesis bounds on the present value of  the  coupling function, $\\alpha_0$, strongly  differ from  some  theories to  others. For  example, in  theories defined by $\\alpha  \\propto \\mid\\varphi\\mid$ analytical   estimates lead to  very stringent nucleosynthesis     bounds    on      $\\alpha_0$  ($\\lesssim 10^{-19}$). By contrast,  in scalar-tensor theories defined by $\\alpha \\propto \\varphi$ much larger limits on $\\alpha_0$ ($\\lesssim 10^{-7}$) are found. ",
        "introduction": "\\label{introduction}  Scalar-tensor (ST) gravity  theories \\cite{B,W,N} have  become a focal point  of interest  in  many   areas  of  gravitational  physics   and cosmology.   They provide the  most natural generalizations of General Relativity (GR) by introducing an additional scalar field, $\\phi$, the dynamical  importance of which is determined  by an arbitrary coupling function $\\omega(\\phi)$.  Indeed,  most   recent attempts  at  unified models  of   fundamental    interactions,  i.e.,   string     theories \\cite{Green88} predict the existence of scalar  partners to the tensor gravity of GR, and would have ST gravity as their low energy limit. In addition, ST theories are  important in cosmology because they provide a natural (non-fine tuned) way of exiting the inflationary epoch.  Solar system experiments  (time delay in  the Viking Mars data,  Lunar Laser Ranging, etc.)  can  put limits on the  present  deviation of ST theories with   respect to GR.  It is   admitted \\cite{Will} that such experiments impose the constraint $\\omega_0>500$  to the present value of the  coupling function  and,  therefore, that  ST theories   are at present very close to   General Relativity.  This limit  on $\\omega_0$ does  not  necessarily imply  that the universe  evolution  is, at any time, very   close    to that  found in   GR.   It  has   been   shown \\cite{DN93,SKW99}   that, in   some  ST   theories,  the  cosmological evolution  drives the scalar   field toward a state  indistinguishable from GR. Within this 'attraction mechanism', the scalar field can play an  important  role only in   early cosmology because,  afterwards, it evolves toward a state with a vanishingly small scalar contribution.  The dynamical importance of the scalar field in the early universe can be checked by means of the observed abundance of light elements, which has to  be explained as  an outcome  of the primordial nucleosynthesis process (PNP).   Contrary to the  weak   field limit, there is  not  a commonly   accepted  PNP  constraint   on $\\omega_0$.     Some authors \\cite{Santiago97,DP99} have recently  found that, in  the framework of some ST  theories, $\\omega_0\\gtrsim10^7$  is  required  to obtain  the observed primordial abundance  of    light elements.  Other    authors \\cite{DTYepes87,SA96b} have instead found  that the PNP test typically imposes the limit $\\omega_0\\gtrsim10^{20}$.  In this   paper we will first  reexamine  the attraction  mechanism of Refs. \\cite{DN93,SKW99} and, then,  we will investigate the reason for the enormous  discrepancy (thirteenth orders of  magnitude) in the PNP bound on $\\omega_0$ obtained by  different authors. We will show that, in general,  the evolution  of the  scalar  field is  governed  by two opposite   mechanisms: an attraction  and  a  repulsion mechanism. The attraction  mechanism  dominates  the recent   epochs  of the universe evolution only if the scalar field and  its derivative satisfy certain boundary  conditions  which  depend on   the  particular scalar-tensor theory   used to describe the  universe  evolution. We will also apply this   generalized formalism  to   the   theories considered   in Ref. \\cite{SA96b},  where  the coupling  function  was  restricted  to be a monotonic function of time. We will show that the nucleosynthesis   bounds  numerically    obtained  in \\cite{SA96b}  ($ \\omega_0\\gtrsim 10^{20} $) are in  close agreement with the analytical estimates for these  theories. When the same  arguments are applied to other  ST theories (as  those of Refs. \\cite{DN93,SKW99}), one obtains much less stringent bounds on $\\omega_0$. Consequently, the particular ST  theory  used to  describe   the universe  evolution has  a crucial importance on the PNP  limit on $\\omega_0$  and, therefore, it is  not possible to establish a general and unique limit for all ST models.  The  plan   of  the paper  is  as   follows. We   begin  outlining the scalar-tensor theories as well as the two frames usually considered to build     up       cosmological   models     in     their    framework (Sec. \\ref{section-ST}). An autonomous    evolution equation for   the scalar field is  then obtained in   Sec. \\ref{Decoupled} both  for the Jordan and the Einstein frame. Using this equation, and the particular family of scalar-tensor theories specified in Sec. \\ref{The-coupling}, we  then  analyze  the  evolution of   the scalar  field both  in  the radiation-dominated  epoch  (Sect.  \\ref{Radiation-epoch})   and   the matter-dominated  epoch  (Sect. \\ref{Matter-epoch}).  Our  results are then applied to estimate  the nucleosynthesis bounds on $\\omega_0$ for this  family   of theories    (Sec. \\ref{nucleosynthesis}).   Finally, conclusions   and     a   summary  of  our      results are   given in Sec. \\ref{conclusions}. ",
        "conclusions": "In  this  paper we  have  analyzed  the  convergence of  scalar-tensor theories (ST)  toward GR and  its consequences on  the nucleosynthesis bounds on the present value of the coupling function.  To that end, we have  deduced an autonomous evolution equation for the Jordan scalar field.   By writing this equation in  Einstein units, we have analyzed the evolution of the scalar  field  both  in  the  radiation-dominated  epoch  and  in  the matter-dominated  epoch.   We  have  considered  a  coupling  function defined  by Eq.   (\\ref{Ecoupling}), which  reproduces all  the models proposed by Barrow \\& Parsons \\cite{BP97} in the limit close to GR. We have then shown that, in general, the evolution of the scalar field is governed  by two opposite  mechanisms: an  attraction and  a repulsion mechanism. The attraction mechanism dominates the recent epochs of the universe  evolution  if,  and  only  if,  the  scalar  field  and  its derivative satisfy  certain boundary  conditions which depend  on each particular scalar-tensor theory.  Our results have been then applied to obtain an analytical estimate of the  Big-Bang nucleosynthesis  (BBN)  bounds on  $\\omega_0$.  We  have found  that the  particular ST  theory used  to describe  the universe evolution   has  a   crucial   importance  on   the   BBN  limits   on $\\omega_0$. Therefore, it  is not possible to establish  a general and unique  limit  for all  ST  models.  In  the  particular  case of  the theories analyzed in this paper, where $\\alpha \\propto |\\varphi|$, our analytical estimates  are in close agreement  with the nucleosynthesis bounds  numerically obtained in  \\cite{SA96b} ($\\alpha^{2}_{0}\\lesssim 10^{-20}$).   Theories different  from  those analyzed  in this  paper could imply very different BBN bounds.  For instance, in the case of a ST theory defined by $\\alpha(\\varphi) \\propto \\varphi$, where only the attraction    mechanism    is   present,    the    BBN   bounds    are ($\\alpha^{2}_{0}\\lesssim  10^{-7}$)  \\cite{SKW99,DP99}.   In the  same way, in  scalar-tensor theories with a non-monotonic  evolution of the speed-up factor,  the BNB  limits are drastically  relaxed to  a value comparable   to   that   obtained   from  solar   system   experiments ($\\alpha^{2}_{0} \\lesssim 0.02$).  In addition, in this last case, the allowed  range for  the  baryon density  is  much wider  than in  the standard GR cosmologies \\cite{SA96b,AS97}.  \\ack{This  work  has  been  partially  supported  by  the Generalitat Valenciana (project number GV00-139-1), Spain.}"
    },
    "0201/astro-ph0201524_arXiv.txt": {
        "abstract": "The linearity and quietness of the Local ($< 10 Mpc$) Hubble Flow (LHF) in view of the very clumpy local universe is a long standing puzzle in standard and in open CDM cosmogony. The question addressed in this paper is whether  the antigravity component of the recently discovered dark energy can cool the velocity flow enough to provide a solution to this puzzle. We calculate the growth of matter fluctuations in a flat universe containing a fraction $\\Omega_X(t_0)$ of dark energy obeying the time independent equation of state $p_X = w \\rho_X$.  We find that dark energy can indeed cool the LHF. However the dark energy parameter values required to make the predicted velocity dispersion consistent with the observed value $v_{rms}\\simeq 40km/sec$ have been ruled out by other observational tests constraining the dark energy parameters $w$ and $\\Omega_X$. Therefore despite the claims of recent qualitative studies dark energy with time independent equation of state can not by itself explain the quietness and linearity of the Local Hubble Flow. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0201/astro-ph0201148_arXiv.txt": {
        "abstract": "We discuss a method to determine orbital properties and masses of low-mass bodies orbiting eclipsing binaries. The analysis combines long-term eclipse timing modulations (light-travel time or LTT effect) with short-term, high-accuracy astrometry. As an illustration of the method, the results of a comprehensive study of Hipparcos astrometry and over a hundred years of eclipse timings of the Algol-type eclipsing binary R Canis Majoris are presented. A simultaneous solution of the astrometry and the LTTs yields an orbital period of $P_{12}=92.8\\pm1.3$~yr, an LTT semiamplitude of $2574\\pm57$~s, an angular semi-major axis of $a_{12}=117\\pm5$ mas, and values of the orbital eccentricity and inclination of $e_{12}=0.49\\pm0.05$, and $i_{12}=91.7\\pm4.7$~deg, respectively. Adopting the total mass of R~CMa of $M_{12}=1.24\\pm0.05$~M$_{\\odot}$, the mass of the third body is $M_3=0.34\\pm0.02$~M$_{\\odot}$ and the semi-major axis of its orbit is $a_3=18.7\\pm1.7$~AU. From its mass, the third body is either a dM3-4 star or, more unlikely, a white dwarf. With the upcoming microarcsecond-level astrometric missions, the technique that we discuss can be successfully applied to detect and characterize long-period planetary-size objects and brown dwarfs around eclipsing binaries. Possibilities for extending the method to pulsating variables or stars with transiting planets are briefly addressed. ",
        "introduction": "Within the next decade several space astrometry missions, FAME and DIVA then SIM and GAIA, capable of sub-milli-arcsecond to micro-arcsecond accuracy are expected to be launched. One of the primary scientific goals of these missions is the astrometric detection of low mass objects around nearby stars, including brown dwarfs and Jupiter-sized planets. The detection of these objects will be accomplished through the observation of the reflex motion of the host star caused by the gravitational pull of the low-mass body. Although these missions are capable of very high astrometric accuracies and can observe up to millions of stars, their lifetimes are relatively short (2.5 to 5 yr). Thus, these space missions are optimized to detect planets within the habitable zones of late-type stars, but they could fail to detect (additional) planets with longer periods. It is important, however, to secure a complete picture of the bodies orbiting a star both from a pure census point of view and also to understand the genesis and evolution of planetary systems. In addition, planets do not necessarily remain within the habitable zone because of long-term chaotic perturbations. As we know from our Solar System, the presence of massive planets, such as Jupiter and Saturn, in distant orbits play a crucial role in stabilizing the orbits of the inner planets.  One effective way of extending the time baseline that permits the discovery of long period exosolar planets or brown dwarfs is to use the light-travel time (hereafter LTT) effect in eclipsing binaries. From this technique, the eclipses act as an accurate clock for detecting subtle variations in the distance to the object (this is analogous to the method used for discovering earth-sized objects around pulsars, see Wolszczan \\& Frail 1992). The periodic quasi-sinusoidal variations of the eclipse arrival times have a very simple and direct physical meaning: the total path that the light has to travel varies periodically as the eclipsing pair moves around the barycenter of the triple system. The amplitude of the variation is proportional both to the mass and to the period of the third body, as well as to the sine of the orbital inclination. As discussed by Demircan (2000), nearly 60 eclipsing binaries show evidence for nearby, unseen tertiary components using LTT effects. A recent example of a brown dwarf detected around the eclipsing binary V471~Tau using this method was presented by Guinan \\& Ribas (2001). Also this method is being employed in selected low mass eclipsing binaries to search for extrasolar planets (Deeg et al. 2000).  The primary advantages of using the LTT effect to detect third bodies in eclipsing binaries are: {\\em i)} The necessary photometry can be secured with small telescopes using photoelectric or CCD detectors; {\\em ii)} The number of eclipsing binaries is large -- 4000 currently known in the Galaxy -- and this number could increase very significantly when results from upcoming astrometry and photometry (e.g., MONS, COROT, Eddington, Kepler) missions are available; {\\em iii)} For select eclipsing binaries (with sharp and deep eclipses) the timings can be determined with accuracies as good as several seconds; {\\em iv)} The mass of the eclipsing pair can be known from conventional spectroscopic and light curve analyses. A shortcoming of the LTT method is that only upper limits to the mass and size of the orbit of the tertiary component can be determined (the analysis yields the mass function\\footnote{$f(M_3)=(M_3^3 \\sin i_3^3)/(M_{12}+M_3)^2$}, $f(M_3)$, and $a_3 \\sin i_3$). However, as it was demonstrated in the case of Algol (Bachmann \\& Hershey 1975), the LTT analysis can be complemented with astrometry to yield the orbital inclination and thus the actual mass and semi-major axis of the third body. Furthermore, with the orbital properties ($P$, $e$ and $\\omega$) known from the LTT analysis, only a small fraction of the astrometric orbit needs to be covered when using high-accuracy astrometry.  In this paper we present the results of the combined LTT analysis and Hipparcos astrometry of the Algol-type eclipsing binary R~CMa. The residuals of over 150 eclipse timings extending from 1887 to 2001 show a periodic ($\\sim$93 yr) quasi-sinusoidal modulation. As previously shown by Radhakrishnan, Sarma, \\& Abhyankar (1984) and Demircan (2000), these variations are best explained by the LTT effect arising from the gravitational influence of a third body. The Hipparcos astrometry also shows the presence of small but significant acceleration terms in the proper motion components explicable by the reflex motion from a third body. Our study illustrates that with a well-defined LTT effect, only a few years of accurate astrometry are needed to constrain the orbital solution and determine the mass of the third body. ",
        "conclusions": "This paper presents a combined analysis of short-term accurate astrometry and long-term timing residuals applied to the eclipsing binary R~CMa. The study yields the complete orbital and physical properties of the tertiary component. A determination of the mass of the third body is possible because the masses of the eclipsing binary components themselves are well-known from light and radial velocity curve analyses.  The example discussed here illustrated the capabilities of a method that will reach its full potential with the upcoming high-accuracy astrometric missions. The improvements in precision of the future astrometric measurements are due to an increase up to a thousand-fold relative to Hipparcos and the quality of the photometry (and thus the eclipse timings) will also improve. More quantitatively, timings with accuracies of $\\sim$10 s are now possible for select eclipsing binaries with sharp eclipses. The detection of large planets ($\\sim$10~M$_{\\rm J}$) in long-period orbits ($\\sim$10--20~yr) around eclipsing binaries will be therefore a relatively easy task. The short-term astrometry will confirm the detections and yield the complete orbital solution (most significantly the inclination) and thus the actual mass of the orbiting body.  One of the unexpected outcomes of the Hipparcos mission has been that a primarily astrometric satellite can also provide valuable new results from its photometric measurements alone (numerous new variables, HD 209458 planetary transits, etc). The data analysis of the next generations of astrometric satellites will surely benefit from a simultaneous analysis of the astrometric and photometric data.  Astrometric missions such as GAIA will likely detect one million new eclipsing binaries (a smaller number is expected for FAME). About one per cent of the eclipsing binaries observed by Hipparcos has a 0.0001 day precision in the reference epoch, which is enough to detect the LTT effect that would arise from a 10 Jupiter mass third body with a 11 year period. If we assume the same ratio for GAIA, hundreds to thousands of third bodies would be detected. Although GAIA astrometry alone will be able to give the orbit for the closest stars, the orbit for more distant stars will depend on the availability of ground-based light curves to define the reference epoch.  This method of combining LTT analysis and astrometry complements very well with the ongoing spectroscopic searches. The LTT analysis favors the detection of long-period third bodies around eclipsing binaries because the amplitude of the time delay due to the LTT effect is proportional to $P_{12}^{2/3}$ while the spectroscopic semi-amplitude is proportional to $P_{12}^{-1/3}$.  When the samples of spectroscopic and LTT systems are sufficiently large, we will have a complete picture of the distribution of bodies in a stellar system and a realistic test of planet formation theories will be possible.  Finally, the LTT analysis method does not have to be necessarily applied to eclipsing binaries. In essence, the method is based upon having a ``beacon in orbit'', which, in the case of eclipsing binaries, are the mid-eclipse times. However, any strictly periodic event that can be predicted with good accuracy could be potentially useful to detect stellar or sub-stellar companions. This includes, for example, pulsating stars. More interestingly, transiting planets are also prime candidates for LTT studies. In this case, not only further orbiting planets could be discovered, but also good chances for detecting moons around the transiting planet exist."
    },
    "0201/astro-ph0201181_arXiv.txt": {
        "abstract": "We have used 850$\\mu$m  maps obtained as part of  the Canada-UK Deep Submillimeter Survey (CUDSS) to investigate the sub-mm properties of  Lyman-break galaxies (LBGs).  We used three samples of Lyman-break galaxies: two  from the Canada-France Deep Fields (CFDF) survey covering CUDSS-14 and CUDSS-3, and one from Steidel and collaborators also covering CUDSS-14. We measure  a mean flux from both CFDF LBG samples at a level of $\\sim$2$\\sigma$ of 0.414 $\\pm$ 0.263  mJy for CUDSS-03 and  0.382 $\\pm$ 0.206 mJy for CUDSS-14, but the Steidel et al. sample is consistent with zero flux.  From this we place upper limits on the Lyman-break contribution to the $850{\\mu}m$ background of  $\\sim$20\\%. We have also measured the cross-clustering between the LBGs and SCUBA sources. From this measurement we infer a large clustering amplitude of $r_o$ = 11.5 $\\pm$ 3.0 $\\pm$ 3.0 $h^{-1}$Mpc for the Steidel et al sample (where the first error is statistical and the second systematic), $r_o$ = 4.5 $\\pm$ 7.0 $\\pm$ 5.0 $h^{-1}$Mpc for CFDF-14 and $r_o$ = 7.5 $\\pm$ 7.0 $\\pm$ 5.0  $h^{-1}$Mpc for CFDF-3. The Steidel et al sample, for which we have most only significant detection of clustering is also the largest of the three samples and has spectroscopically  confirmed redshifts ",
        "introduction": "Recent work at optical and submillimeter (sub-mm) wavelengths has granted us unprecedented access to the high-redshift universe.  In particular, the Lyman-break selection technique \\citep{ste93} has provided thousands of star forming galaxies at redshifts z$\\sim$3 and z$\\sim$4.  These data, and the results of other optical surveys \\citep{lil96,mad98,hog98}, have shown that the global star formation rate  (inferred from the $UV$ luminosity  density at different redshifts),  increases with redshift to  z$\\sim$1 and may remain constant to at least z$\\sim$4 \\citep{saw97,ste99}, implying the beginning of the epoch of galaxy formation occurred at $z>>$ 1.  \\  Studies of the spatial distribution of Lyman-break galaxies  show that they  are highly clustered even at these early redshifts \\citep{gia98}.  This was initially unexpected since the clustering of galaxies had been shown to decrease with redshift to z$\\sim$1 \\citep{carl97, lefev96} in line with theoretical predictions, but the strong clustering of LBGs is  actually a natural consequence of the effects of bias. \\citet{kai84} showed that the high peaks of the density distribution in the  early universe will have been highly clustered, and so objects that form from these high peaks, clusters at a redshift of zero or galaxies at a  redshift  of 3, should also be highly clustered.  \\  The high star formation rates (50-100 M$_{\\odot}$ yr$^{-1}$) and comoving density of the Lyman-break galaxies make them attractive progenitors of present-day elliptical galaxies \\citep{pet98} though their masses are still highly uncertain \\citep{saw98,som01}. However, the Submillimeter Common-User Bolometer Array (SCUBA) on the James Clerk Maxwell Telescope (JCMT) has revealed a population of dusty galaxies with implied star formation rates of $>$ 300 M$_{\\odot}$ yr$^{-1}$ \\citep{sma97,hug98,eal99}.  The redshifts of these  objects are still  highly uncertain, with estimates of the median redshift  of the population lying between 2 and 3 \\citep{eal00,bar99,sma00,yun02,fox02}.  They have  similar spectral energy distributions  to  today's ultraluminous infrared  galaxies (ULIRGs) and often show disturbed morphology or multiple  components,  implying they may be the result of galaxy mergers \\citep{lil99,ivi00}.   Both LBGs and SCUBA sources have sufficient star formation rates to form  present-day elliptical galaxies but the extremely high star formation rates of the latter mean this can be done on the order of 10$^8$ years, as the homogeneous properties  of local ellipticals indicate is the case.  \\   The nature of the relationship between these two populations remains unclear. An obvious scenario is one in which they form a continuum of objects with the bright sub-mm selected sources representing the highest star forming Lyman-break galaxies.   \\citet{adel00} claim that by assuming an $L_{bol,dust}/L_{UV}$ typical of normal starbursts, the LBG population can produce the  bulk  of the  850 $\\mu$m background. In this picture the two populations are the same objects and a separate population of  highly obscured ULIRG-like objects is not needed.    The ratio of optical to sub-mm emission for the brighter SCUBA sources is, however, much less than for the starbursts considered by Adelberger \\& Steidel and so these objects almost certainly represent a separate population \\citep{gea00, eal00,dow99}, but it is unclear whether the fainter SCUBA sources, with S$_{850} <$ 3 mJy, overlap with the LBG population. \\   Various optical techniques have been used to infer the dust content of LBG's. \\citet{pet01} have used optical line ratios  and \\citet{sha01} fitted the predictions of star-formation models to  optical and near-IR photometry. Both have concluded that the most intrinsically luminous LBGs, which have higher star formation rates, contain more dust.  However,  a more reliable way of measuring the dust content is directly through sub-mm photometry. \\citet{cha00} have used SCUBA to observe a sample of  high-SFR LBGs and they estimate that the 850$\\mu$m flux density is at least two times lower than predicted from $UV$ colours. Using the sub-mm map of the Hubble Deep Field (HDF), \\citet{pea00} statistically detected the sub-mm flux of galaxies with high $UV$ luminosities, and thus high star formation rates.  They detect a higher mean flux of 0.2 $\\pm$0.04 mJy  (for galaxies with an inferred star formation rate (SFR) of 1 $h^{-2}$M$_{\\odot}$ yr$^{-1}$).   \\    This paper examines  the relationship between Lyman-break galaxies and SCUBA sources in  two Canada-UK Deep Submillimeter Survey (CUDSS) fields and is organised as follows.  \\S 2 describes the sub-mm and optical/$UV$ data.  In \\S 3 we investigate the sub-mm fluxes of  LBGs.   \\S 4 discusses the dust properties of  LBGs that can be inferred from the results in \\S 3. In \\S 5 the correlation function between the two populations is presented.  In  \\S 6 we discuss the results and their implications.\\ ",
        "conclusions": "We have used the 850 $\\mu$m maps from the Canada-UK Deep Sub-mm Survey to study (i) the sub-mm flux and dust properties of Lyman-break galaxies and  (ii) the angular correlation between Lyman-break galaxies and SCUBA sources.  We obtain the following results: \\  \\  1. We marginally detect (at the 2$\\sigma$ level) sub-mm flux from Lyman-break galaxies in the CFDF-14 and CFDF-03 samples but we do not detect flux from the Steidel et al sample. The flux levels are: 0.382 $\\pm$ 0.206 mJy for the 14-hour field and 0.414 $\\pm$ 0.263 mJy for the 3-hour field.  Further, we show that possibly because of LBG-SCUBA clustering SCUBA sources not identified with a LBG galaxy must be removed from the map before a proper analysis can be performed. \\  2. Lyman-break galaxies are the best optical identification for four SCUBA sources although it is possible that some of these identifications may be incorrect.  There are indications that these objects may lie in a region of spatial over-density.   \\  3.  An upper limit for the dust mass of Lyman-break galaxies was calculated from their sub-mm flux results and we conclude that these masses can be no larger that those of near-by galaxies. \\  4.  The SCUBA-LBG correlation function was measured for all three sample of Lyman-break galaxies. We found a high-amplitude $r_o$= 11.5 $\\pm$ 3.0 $\\pm$ 3.0  $h^{-1}$ Mpc for the Steidel et al sample, $r_o$ = 4.5 $\\pm$ 7.0 $\\pm$ 5.0 $h^{-1}$ Mpc for CFDF-14 and $r_o$ = 7.5 $\\pm$ 7.0 $\\pm$ 5.0  $h^{-1}$ Mpc for CFDF-03, (where the first error is statistical and the second systematic).  \\  {\\it Acknowledgments} We are grateful to the many members of the staff of the Joint Astronomy Centre who have helped us with this project.  Research by Simon Lilly is supported by the National Sciences and Engineering Council of Canada and by the Canadian Institute of Advanced Research. Research by Tracy Webb is supported by the National Sciences and Engineering Council of Canada,the Canadian National Research Council and the Ontario Graduate Scholarship Program.  Research by Stephen Eales, David Clements, Loretta Dunne and Walter Gear is supported by the Particle Physics and Astronomy Research Council.  Stephen Eales also acknowledges support from Leverhulme Trust. The JCMT is operated by the Joint Astronomy Centre on behalf of the UK Particle Physics and Astronomy Research Council, the Netherlands Organization for Scientific Research and the Canadian National Research Council.   We also thank Ray Carlberg for many helpful discussions.   \\"
    },
    "0201/astro-ph0201304_arXiv.txt": {
        "abstract": "{ This is the second paper in a series of three in which we present an exhaustive inventory of the solid state emission bands observed in a sample of 17 oxygen-rich dust shells surrounding evolved stars. The data were taken with the Short and Long Wavelength Spectrographs on board of the Infrared Space Observatory (ISO) and cover the 2 to 200 $\\mu$m wavelength range. Apart from the broad 10 and 18 $\\mu$m bands that can be attributed to amorphous silicates, at least 49 narrow bands are found whose position and width indicate they can be attributed to crystalline silicates. Most of these emission bands are concentrated in well defined spectral regions (called complexes). We define 7 of these complexes; the 10, 18, 23, 28, 33, 40 and 60 micron complex. We derive average properties of the individual bands. Almost all of these bands were not known before ISO. Comparison with laboratory data suggests that both olivines (Mg$_{2x}$Fe$_{(2-2x)}$SiO$_4$) and pyroxenes (Mg$_x$Fe$_{(1-x)}$SiO$_3$) are present, with x close to 1, i.e. the minerals are very Mg-rich and Fe-poor. This composition is similar to that seen in disks surrounding young stars and in the solar system comet Hale-Bopp. A significant fraction of the emission bands cannot be identified with either olivines or pyroxenes. Possible other materials that may be the carriers of these unidentified bands are briefly discussed. There is a natural division into objects that show a disk-like geometry (strong crystalline silicate bands), and objects whose dust shell is characteristic of an outflow (weak crystalline silicate bands). In particular, stars with the 33.5~$\\mu$m olivine band stronger than about 25 percent over continuum are invariably disk sources. Likewise, the 60 $\\mu$m region is dominated by crystalline silicates in the disk sources, while it is dominated by crystalline H$_{2}$O ice in the outflow sources. We show that the disk and outflow sources have significant differences in the shape of the emission bands. This difference must be related to the composition or grain shapes of the dust particles. The incredible richness of the crystalline silicate spectra observed by ISO allows detailed studies of the mineralogy of these dust shells, and is the origin and history of the dust.} ",
        "introduction": "\\label{sec:obser}  At the end of their life both low and high mass stars loose a large fraction of their mass in the form of a dense stellar wind. When the temperature in the outer regions of the atmosphere becomes low enough solid material ({\\em dust}) condenses out of the gas. Mass loss can eventually dominate stellar evolution. Since the dust may also play an important role in the mass loss process, it is interesting to investigate the physical and chemical processes responsible for dust formation. One way to do this is to study the endproducts of this dust formation process. The composition of the dust which has condensed provides valuable information on the conditions when the dust was formed and thereafter.  The dust around stars can be observed at different wavelengths. In the visible and near-infrared~(NIR) one can look at the wavelength dependence of the absorption caused by the dust. Another way to investigate the dust at these wavelengths is by means of scattered light of the central source by small dust particles in the circumstellar environment or with polarimetric observations. In this paper we will study the emission of the dust in the mid-infrared (MIR) and far-infrared~(FIR).  The dusty environments around evolved stars can be divided in carbon-rich (C-rich) and oxygen-rich (O-rich) environments, depending on the C/O ratio of the mass-losing stars. This division is the result of the stability of the CO molecule, which is formed before the dust condenses. If there is more carbon than oxygen (C/O $>1$) all the oxygen will be trapped in CO, and the dust species will be carbonaceous, e.g. SiC, Polycyclic Aromatic Hydrocarbons (PAHs) or amorphous carbon. If the C/O ratio is smaller than one, i.e. there is more oxygen than carbon, all the carbon is trapped in CO and O-rich dust will be formed, e.g. simple oxides and silicates.  In 1995 the Infrared Space Observatory (ISO; Kessler et al. 1996) was launched which opened the possibility to study mass-losing stars at infrared wavelengths with unprecedented wavelength coverage and spectral resolution. Before the launch of ISO it was generally assumed that in the dusty winds of O-rich evolved stars only amorphous silicates were formed. One of the remarkable discoveries of ISO was the detection of crystalline silicates outside our own Solar System. Their spectral signature was not only found in the spectra of young stars (Waelkens et al, 1996), but also in the outflows of evolved stars (Waters et al, 1996). In the latter case, we even found one example, IRAS09425-6040, where 75\\% of the circumstellar dust consists of crystalline silicates (Molster et al. 2001).  In contrast to amorphous silicates, crystalline silicates provide a unique opportunity to determine for the first time the chemical composition of the dust particles. The relatively sharp features of the crystalline silicates are very sensitive to compositional changes, this in contrast to the broad and smooth amorphous silicate features. Crystalline silicates will also help us to better understand the physical and chemical conditions under which dust is formed. In particularly, they need high temperatures and a slow cooling down to form.  Crystalline silicates are found in the outflows of evolved stars which replenish the ISM, and also around young stars which form from the ISM (e.g. Waters et al. 1996; Waelkens et al. 1996). However, up to now no crystalline silicates have been found in the ISM. The abundance in the ISM might be too low. However, this would make it difficult to explain the high abundance of crystalline silicates in a young star as HD100546 (Waelkens et al. 1996; Malfait et al. 1998). On the other hand, if the crystalline silicates are destroyed in the ISM then the question arises how they are formed around young stars. Possible solutions to this last problem are given by Molster et al. (1999a; hereafter MYW).  The crystallization process is not well understood and in order to better constrain the origin of these grains, it is important to accurately describe the observed properties in different environments. After the first reports of the crystalline silicates, only a few objects have been analysed in some detail, e.g. HD100546 (Malfait et al. 1998), CPD$-56^{\\circ}8032$ (Cohen et al. 1999), AFGL4106 (Molster et al. 1999b), IRAS09425-6040 (Molster et al. 2001). This will be the first attempt to obtain an average spectrum of the different crystalline silicate features which are found around evolved stars.  In the first paper of this series (Molster et al. submitted; hereafter paper I) we have described the infrared spectra of 17 sources. Here, we will derive mean spectra based on these sources and identify the different features. In the third paper of this series (Molster et al. submitted; hereafter paper III) we will apply a simple dust emission model to the spectra of our sample and derive several (spectral) trends.  In Section~\\ref{sec:obser} we will briefly summarise how we obtained these mean spectra. For a more extensive description we refer to paper I. The results will be shown in Section~\\ref{sec:results}, where we will also order and define the different solid state bands and complexes found. In Section~\\ref{sec8:ident} we will identify most of the features found in Section~\\ref{sec:results} by comparison to laboratory spectra of cosmic dust analogues.  \\begin{table*}[b] \\caption[]{The characteristics of the different features. The mean $\\lambda$ (FWHM) is the average wavelength (Full Width Half Max) of the feature, with the errors taken into account. $\\lambda$ (FWHM) min and max are the minimum and maximum value found in the sample. The N indicates not a member of one of the 7 complexes} {\\small \\begin{tabular}{|lrr|rrr|l|c|} \\hline \\multicolumn{3}{|c|}{$\\lambda$}& \\multicolumn{3}{|c|}{FWHM}& Identification & complex\\\\ mean & min & max & mean & min & max & & nr\\\\ \\hline \\multicolumn{1}{|c}{8.3} &     8.20 &     8.40 &     .42 &      .41 &      .43 & & 10 \\\\ \\multicolumn{1}{|r}{9.14} &     9.12 &     9.17 &     .30 &      .24 &      .68 & silica? & 10 \\\\ \\multicolumn{1}{|r}{9.45} &     9.45 &     9.46 &     .19 &      .15 &      .25 & & 10 \\\\ \\multicolumn{1}{|c}{9.8} &     9.77 &     9.84 &     .17 &      .14 &      .29 & forsterite + enstatite & 10 \\\\ 10.1 &     9.59 &    10.61 &    2.56 &     1.30 &     3.77 & amorphous silicate& 10 \\\\ 10.7 &    10.57 &    10.90 &     .28 &      .11 &      .66 & enstatite & 10 \\\\ 11.05 &    11.04 &    11.06 &     .05 &      .03 &      .11 & instrumental artifact & 10 \\\\ 11.4 &    11.33 &    11.50 &     .48 &      .38 &      .86 & forsterite, diopside? & 10 \\\\ % \\hline 15.2 &    15.00 &    15.42 &     .26 &      .13 &      .73 &  enstatite & 18 \\\\ 15.9 &    15.69 &    16.06 &     .43 &      .24 &      .65 &  silica?? & 18 \\\\ 16.2 &    16.10 &    16.37 &     .16 &      .08 &      .62 &  forsterite & 18 \\\\ \\multicolumn{1}{|l}{16.50}&    16.49 &    16.50 &     .11 &      .11 &      .11 &  PAH?& 18 \\\\ 16.9 &    16.73 &    17.07 &     .57 &      .37 &      .84 &  a blend of 16.7 and 17.0 & 18 \\\\ 17.5 &    16.79 &    18.46 &    2.10 &      .81 &     3.66 &  amorphous silicate      & 18 \\\\ 17.5 &    17.43 &    17.61 &     .18 &      .13 &      .36 &  enstatite & 18 \\\\ 18.0 &    17.90 &    18.16 &     .48 &      .28 &     1.24 &  enstatite + forsterite & 18 \\\\ 18.9 &    18.43 &    19.17 &     .62 &      .36 &     1.20 &  forsterite? & 18 \\\\ 19.5 &    19.36 &    19.75 &     .40 &      .14 &      .86 &  forsterite + enstatite & 18 \\\\ \\hline 22.4 &    22.26 &    22.51 &     .28 &      .16 &      .55 &   & 23 \\\\ 23.0 &    22.82 &    23.14 &     .48 &      .28 &      .72 &  enstatite & 23 \\\\ 23.7 &    23.45 &    23.81 &     .79 &      .54 &     1.29 &  forsterite & 23 \\\\ 23.89&    23.88 &    23.90 &     .18 &      .13 &      .25 &   & 23 \\\\ 24.5 &    24.16 &    24.65 &     .42 &      .16 &     1.04 &  enstatite + ? & 23 \\\\ 25.0 &    24.83 &    25.14 &     .32 &      .25 &      .53 &  diopside? & 23 \\\\ \\hline 26.8 &    26.71 &    26.93 &     .37 &      .21 &      .47 &   & 28 \\\\ 27.6 &    27.46 &    27.79 &     .49 &      .28 &     1.18 &  forsterite & 28 \\\\ 28.2 &    27.97 &    28.45 &     .42 &      .23 &      .90 &  enstatite   & 28 \\\\ 28.8 &    28.68 &    28.88 &     .24 &      .19 &      .42 &   & 28 \\\\ 29.6 &    29.37 &    29.90 &     .89 &      .58 &     1.99 &  2 features?, diopside? & 28 \\\\ 30.6 &    30.48 &    30.77 &     .32 &      .18 &      .81 &   & 28 \\\\ 31.2 &    31.12 &    31.27 &     .24 &      .21 &      .36 &  forsterite? & 28 \\\\ \\hline 32.2 &    32.06 &    32.51 &     .46 &      .24 &      .75 &  diopside? & 33 \\\\ 32.8 &    32.56 &    33.03 &     .60 &      .36 &     1.00 &   & 33 \\\\ 32.97&    32.96 &    32.99 &     .20 &      .11 &      .28 &  instrumental artifact & 33 \\\\ 33.6 &    33.45 &    33.71 &     .70 &      .52 &     1.15 &  forsterite & 33 \\\\ 34.1 &    33.93 &    34.36 &     .36 &      .17 &      .74 &  enstatite + diopside? & 33 \\\\ 34.9 &    34.67 &    35.35 &    1.36 &      .63 &     1.88 &  clino-enstatite? & 33 \\\\ 35.9 &    35.76 &    36.20 &     .53 &      .37 &      .88 &  ortho-enstatite? & 33 \\\\ 36.5 &    36.44 &    36.72 &     .39 &      .25 &      .97 &  forsterite + ? & 33 \\\\ \\hline 39.8 &    39.44 &    40.36 &     .74 &      .21 &     2.57 &  diopside? & 40 \\\\ 40.5 &    40.34 &    40.80 &     .93 &      .53 &     1.53 &  enstatite (a blend?) & 40 \\\\ 41.8 &    41.52 &    42.11 &     .72 &      .42 &     1.73 &  41 micron plateau & 40 \\\\ 43.0 &    42.55 &    43.07 &     .89 &      .51 &     1.59 &  cryst. H$_2$O-ice + clino-enst. & 40 \\\\ 43.8 &    43.30 &    44.05 &     .78 &      .41 &     3.01 &  ortho-enstatite & 40 \\\\ 44.7 &    44.39 &    45.13 &     .58 &      .42 &     1.16 &  clino-enstatite, diopside? & 40 \\\\ \\hline 52.9 &    51.40 &    56.55 &    3.11 &     1.75 &     5.91 &  cryst. H$_2$O-ice & 60 \\\\ 62.  &    61.24 &    65.56 &   11.80 &     4.67 &    15.29 &  cryst. H$_2$O-ice (62 micron) + & 60 \\\\ &          &          &         &          &          &  enstatite?, diopside? (65 micron) & 60 \\\\ 69.0 &    68.79 &    69.15 &     .63 &      .46 &     1.04 &  forsterite & 60 \\\\ \\hline 13.5 &    13.40 &    13.60 &     .25 &      .17 &      .44 & responsivity  & N \\\\ 13.8 &    13.74 &    13.90 &     .20 &      .16 &      .23 & enstatite, responsivity?  & N \\\\ 14.2 &    14.15 &    14.28 &     .28 &      .18 &      .55 & enstatite, responsivity?  & N \\\\ 20.7 &    20.54 &    20.84 &     .31 &      .16 &      .84 & silica?, diopside? & N \\\\ 21.5 &    21.35 &    21.65 &     .35 &      .15 &      .79 &  & N \\\\ 26.1 &    25.91 &    26.29 &     .57 &      .22 &     1.10 & forsterite + silica? & N \\\\ 38.1 &    37.83 &    38.21 &     .57 &      .56 &      .61 &  & N \\\\ 47.7 &    47.55 &    47.83 &     .97 &      .78 &     1.85 & FeSi?, a silicate  & N \\\\ 48.6 &    48.34 &    49.07 &     .61 &      .44 &     1.37 & a silicate  & N \\\\ 90.9 &    89.59 &    91.12 &   14.44 &    12.72 &    17.63 &  & N \\\\ \\hline \\end{tabular}} \\label{tab:complex} \\end{table*}  \\begin{table*} \\caption[]{The presence of the features which are seen in at least three different SWS spectra of our sample, + means emission, - is absorption, o means no good data available, 2 implies that a feature split in 2 separate features in this source and underlined x's mean all present, but measured as one feature. For MWC922 the blend of the 9.1 and 9.5 and for OH26.5+0.6 the blend of the 23.0 and 23.7 micron features are in absorption, while all other blends are in emission. The order of the stars is in decreasing strength of the blend of the 33.6 micron feature. am = amorphous silicate, d= diopside; (o/c)e = (ortho/clino)enstatite; f = forsterite; i = instrumental artifact; s= silica; si = a crystalline silicate; w = crystalline water-ice. } {\\normalsize \\begin{tabular}{|l|c@{\\ \\ }c@{\\ \\ }c@{\\ \\ }c@{\\ \\ } c@{\\ \\ }c@{\\ \\ }c@{\\ \\ }c@{\\ \\ }c@{\\ \\ }c@{\\ \\ }c@{\\ \\ } c@{\\ \\ }c@{\\ \\ }c@{\\ \\ }c@{\\ \\ }c@{\\ \\ }c@{\\ \\ }c@{\\ \\ } c@{\\ \\ }c@{\\ \\ }c@{\\ \\ }c@{\\ \\ }c@{\\ \\ }c@{\\ \\ }c@{\\ \\ } c@{\\ \\ }|} \\hline & & & & & & & & & & & & & & & & & & & & & & & & & &  \\\\ & & & & & & & & & & & & & & & & & & & & & & & & & &  \\\\ & & & & & & & & & & & & & & & & & & & & & & & & & &  \\\\ Bandname & { \\begin{rotate}{90} 8.3   \\end{rotate}  }& % { \\begin{rotate}{90} 9.1 s?  \\end{rotate}  }& % { \\begin{rotate}{90} 9.5   \\end{rotate}  }& % { \\begin{rotate}{90} 9.8 f,e  \\end{rotate}  }& % { \\begin{rotate}{90} 10 (am.) \\end{rotate} }& % { \\begin{rotate}{90} 10.7 e \\end{rotate}  }& % { \\begin{rotate}{90} 11.4 f,d? \\end{rotate}  }& % { \\begin{rotate}{90} 15.2 e \\end{rotate}  }&     % { \\begin{rotate}{90} 15.9 s? \\end{rotate}  }&     % { \\begin{rotate}{90} 16.2 f \\end{rotate}  }&     % { \\begin{rotate}{90} 16.9  \\end{rotate}  }&     % { \\begin{rotate}{90} 17.5 e \\end{rotate}  }&     % { \\begin{rotate}{90} 18 (am.) \\end{rotate} }& % { \\begin{rotate}{90} 18.0 f,e \\end{rotate}  }&     % { \\begin{rotate}{90} 18.9 f? \\end{rotate}  }&     % { \\begin{rotate}{90} 19.5 f,e \\end{rotate}  }&     % { \\begin{rotate}{90} 20.7 s?,d? \\end{rotate}  }&     % { \\begin{rotate}{90} 21.5  \\end{rotate}  }&     % { \\begin{rotate}{90} 22.4  \\end{rotate}  }&     % { \\begin{rotate}{90} 23.0 e \\end{rotate}  }&     % { \\begin{rotate}{90} 23.7 f \\end{rotate}  }&     % { \\begin{rotate}{90} 23.89 \\end{rotate}  }&     % { \\begin{rotate}{90} 24.5 e \\end{rotate}  }& % { \\begin{rotate}{90} 25.0 d? \\end{rotate}  }& % { \\begin{rotate}{90} 26.1 f,s? \\end{rotate}  }& % { \\begin{rotate}{90} 26.8  \\end{rotate}  }\\\\    % \\hline IRAS09425       & & & & & &?& &+&+&+&+& & &+&+&+&+&+&+&+&+& &+&+&+&+ \\\\%09 IRAS09425 NGC6537         & & & & & & & & & & &+& & &+&+&+&+&+& &+&+& &+&+&+&+ \\\\%13 NGC6537 NGC6302         & & & & & & & & & & & & & &+&+&+&+&+&+&+&+& &+&+& &  \\\\%01 NGC6302 MWC922          & &\\multicolumn{2}{@{}c@{\\ \\ }}{\\underline{x~~~x}}&-& &-& &+& & &2&+&-&+&+&+&+&+& &+&+&+&+&+&+&+ \\\\%03 MWC922 AC Her          &+&+&+&+&+&+&+&+&+&+&+&+&?&+&+&+&+&+&+&\\multicolumn{2}{@{}c@{\\ \\ }}{\\underline{x~~~x}}&+&+&+& &  \\\\%04 AC Her HD45677         &+&+&+&+&+&+&+&+&+& &+&+&+&+&+&+&+&+&+&+&+& &+&+&+&+ \\\\%16 HD45677 89 Her          &+&+&+&+&+&+&+&+&+&+&+& &+&+&+&+& & & &+&+& &\\multicolumn{2}{@{}c@{\\ \\ }}{\\underline{x~~~x}}& &  \\\\%08 89 Her MWC300          & & & & &-& & & & & &+&+&-&+&+&+&+& &+&+&+& &+&+&+&  \\\\%10 MWC300 VY 2-2          & & & & &+& & & & & & & & & & & &+&+& &\\multicolumn{3}{@{}c@{\\ \\ }}{\\underline{x~~~x~~~x}}&\\multicolumn{2}{@{}c@{\\ \\ }}{\\underline{x~~~x}}& &  \\\\%14 VY 2-2 HD44179         & & & & & & & & &+&+&+& &?&+&+&+&+&+& &+&+& &+&+&+&  \\\\%02 HD44179 HD161796        & & & & &+& & &+&+&+& & &+&+&+&+&+&+& &+&+& &+&+&+&  \\\\%06 HD161796 OH26.5+0.6      & & & & &-& & &-&-&-&-&-&-&-&-&-&-& &-&\\multicolumn{2}{@{}c@{\\ \\ }}{\\underline{x~~~x}}&-&-&-&-&- \\\\%12 OH26.5+0.6 Roberts 22      & & & & & & & & &\\multicolumn{2}{@{}c@{\\ \\ }}{\\underline{x~~~x}}&?&?&+&?&?&?&+&+& &\\multicolumn{5}{@{}c@{\\ \\ }}{\\underline{x~~~x~~~x~~~x~~~x}}& &o \\\\%15 Roberts 22 HD179821        & & & & &+& & & &\\multicolumn{2}{@{}c@{\\ \\ }}{\\underline{x~~~x}}&+& &+&\\multicolumn{2}{@{}c@{\\ \\ }}{\\underline{x~~~x}}&+&+& &+&+&+& &\\multicolumn{2}{@{}c@{\\ \\ }}{\\underline{x~~~x}}&+&  \\\\%07 HD179821 AFGL4106        & & & & &+& & & &\\multicolumn{2}{@{}c@{\\ \\ }}{\\underline{x~~~x}}&+& &+&+&+&+&+&+&+&+&+& &+&+&+&  \\\\%05 AFGL4106 NML Cyg         & & & & &-& & &+&+&+&+& &?&+&+& &+&+& &+&+& &+& & &o \\\\%11 NML Cyg IRC+10420       & &+&+& &+&?&+&+&+&+&+& &+&\\multicolumn{3}{@{}c@{\\ \\ }}{\\underline{x~~~x~~~x}}&+&+& &+&+& &+& &+&  \\\\%17 IRC+10420 \\hline & & & & & & & & & & & & & & & & & & & & & & & & & &  \\\\ & & & & & & & & & & & & & & & & & & & & & & & & & &  \\\\ & & & & & & & & & & & & & & & & & & & & & & & & & &  \\\\ Bandname & { \\begin{rotate}{90} 27.6 f \\end{rotate}  }& % { \\begin{rotate}{90} 28.2 e \\end{rotate}  }& % { \\begin{rotate}{90} 28.8  \\end{rotate}  }& % { \\begin{rotate}{90} 29.6 d? \\end{rotate}  }& % { \\begin{rotate}{90} 30.6  \\end{rotate}  }& % { \\begin{rotate}{90} 31.2 f? \\end{rotate}  }& % { \\begin{rotate}{90} 32.2 d? \\end{rotate}  }& % { \\begin{rotate}{90} 32.8  \\end{rotate}  }&     % { \\begin{rotate}{90} 32.97 i \\end{rotate}  }&     % { \\begin{rotate}{90} 33.6  f \\end{rotate}  }& % { \\begin{rotate}{90} 34.1  e,d? \\end{rotate}  }&     % { \\begin{rotate}{90} 34.9  ce? \\end{rotate}  }&     % { \\begin{rotate}{90} 35.9  oe? \\end{rotate}  }&     % { \\begin{rotate}{90} 36.5  f \\end{rotate}  }& % { \\begin{rotate}{90} 39.8  d? \\end{rotate}  }& % { \\begin{rotate}{90} 40.5  e \\end{rotate}  }& % { \\begin{rotate}{90} 41.8  \\end{rotate}  }& % { \\begin{rotate}{90} 43.0  w,ce \\end{rotate}  }& % { \\begin{rotate}{90} 43.8  oe \\end{rotate}  }& % { \\begin{rotate}{90} 44.7  ce,d?\\end{rotate}  }& % { \\begin{rotate}{90} 47.7  si \\end{rotate}  }&    % { \\begin{rotate}{90} 48.6  si \\end{rotate}  }&    % { \\begin{rotate}{90} 52    w \\end{rotate}  }&    % { \\begin{rotate}{90} 62    w,d \\end{rotate}  }&    % { \\begin{rotate}{90} 69.0  f \\end{rotate}  }&    % { \\begin{rotate}{90} 91     \\end{rotate}  }\\\\   % \\hline IRAS09425       &+&+& &+&+& &+&+& &+&+&+&+&+&+&+&+&+&+&+& & & & & & \\\\%09 IRAS09425 NGC6537         &\\multicolumn{2}{@{\\ \\ }c@{\\ \\ }}{\\underline{x~~~x}}&+&+&+& &+&+& &+&+&\\multicolumn{2}{@{}c@{\\ \\ }}{\\underline{x~~~x}}&+&+&+&+&+&+&+&+&+&+&+&+&+\\\\%13 NGC6537 NGC6302         &+&+& &+&+&+&+&+& &+&+&+&+&+&+&+&+&+&+&+&+&+&+&+&+&+\\\\%01 NGC6302 MWC922          &+&+&+&+&+&+&+&+& &+&\\multicolumn{2}{@{}c@{\\ \\ }}{\\underline{x~~~x}}&+&+& &+&+&+&+&+&\\multicolumn{2}{@{}c@{\\ \\ }}{\\underline{x~~~x}}&+&+&+& \\\\%03 MWC922 AC Her          &+&+& &+&+&?&+&+& &\\multicolumn{2}{@{}c@{\\ \\ }}{\\underline{x~~~x}}& &\\multicolumn{2}{@{}c@{\\ \\ }}{\\underline{x~~~x}}&\\multicolumn{2}{@{}c@{\\ \\ }}{\\underline{x~~~x}}& & &+& & & & & & & \\\\%04 AC Her HD45677         &+&+& &+&+& &+&+& &+&\\multicolumn{2}{@{}c@{\\ \\ }}{\\underline{x~~~x}}&\\multicolumn{2}{@{}c@{\\ \\ }}{\\underline{x~~~x}}&+&+&+&+&+&+& & & & & & \\\\%16 HD45677 89 Her          &+&+& &?& & & &+& &+&\\multicolumn{2}{@{}c@{\\ \\ }}{\\underline{x~~~x}}& & & & & & & & & & & & & & \\\\%08 89 Her MWC300          &\\multicolumn{2}{@{\\ \\ }c@{\\ \\ }}{\\underline{x~~~x}}&+&+&+&+& &+& &+&+&\\multicolumn{3}{@{}c@{\\ \\ }}{\\underline{x~~~x~~~x}}& &+& &\\multicolumn{3}{@{}c@{\\ \\ }}{\\underline{x~~~x~~~x}}& & & & & & \\\\%10 MWC300 VY 2-2          & & & &+&+& & &+& &+&+&+&+&+& &+&\\multicolumn{3}{@{}c@{\\ \\ }}{\\underline{x~~~x~~~x}}& & & &+&+& &+\\\\%14 VY 2-2 HD44179         &+&+& &+&+& &+&+& &+&+&+&+&+& &+& & &+& &\\multicolumn{2}{@{}c@{\\ \\ }}{\\underline{x~~~x}}&+&+&+& \\\\%02 HD44179 HD161796        &+&+&?&+&+& &+&+& &+&+&+&\\multicolumn{2}{@{}c@{\\ \\ }}{\\underline{x~~~x}}&+&+&+&+&?& &\\multicolumn{2}{@{}c@{\\ \\ }}{\\underline{x~~~x}}&+&2&+& \\\\%06 HD161796 OH26.5+0.6      & & & &+& & & &\\multicolumn{4}{@{}c@{\\ \\ }}{\\underline{x~~~x~~~x~~~x}}&+&+&+&+&+&+&+&+&+&\\multicolumn{2}{@{}c@{\\ \\ }}{\\underline{x~~~x}}&+&+&+& \\\\%12 OH26.5+0.6 Roberts 22      &+&+&+&o& & & &+& &+&+&+&\\multicolumn{2}{@{}c@{\\ \\ }}{\\underline{x~~~x}}&\\multicolumn{2}{@{}c@{\\ \\ }}{\\underline{x~~~x}}&+&+&+& & & & &+&?&+\\\\%15 Roberts 22 HD179821        &+&+& &+& & &+&+& &+&+&+&+& &+&+&+&+&+&+&+&+&+&+&+& \\\\%07 HD179821 AFGL4106        &+&+& &o&o& &+&+&+&+&+&+&+&+&+&+&+&+&+&+&\\multicolumn{2}{@{}c@{\\ \\ }}{\\underline{x~~~x}}& &+& & \\\\%05 AFGL4106 NML Cyg         &o&o&o&o&+& &+&+&+&+&+&+&+& &\\multicolumn{2}{@{}c@{\\ \\ }}{\\underline{x~~~x}}&+&+&+&+&\\multicolumn{2}{@{}c@{\\ \\ }}{\\underline{x~~~x}}& &+&+& \\\\%11 NML Cyg IRC+10420       &+&+& &o& & & &+&+&+& &+&+& &?&+&+&+& &+&+&+&+&+& & \\\\%17 IRC+10420 \\hline \\end{tabular}} \\label{tab:overview2} \\end{table*} ",
        "conclusions": "In this paper we have combined the spectra of 17~stars with oxygen-rich circumstellar dust environments. These combined spectra show a wealth of dust features, which now are studied in a systematic way. Most of the features cluster into one of seven complexes (at 10, 18, 23, 28, 33, 40 or 60~$\\mu$m). Each of these complexes is a convolution of a limited set of components, which vary independently.  Our conclusion from this work is that there is a clear distinction between the disk and outflow sources. Not only in the strength  of the features and therefore in abundance of the crystalline silicates, which was the original natural separator profile between these two groups of objects, but also in the characteristics of the complexes.\\\\ \\\\ \\noindent{\\em 10 + 18 micron complex:}\\\\ The difference in shape, a smooth profile for the outflow sources and narrow features for the disk sources, is mainly due to an abundance difference between the crystalline and amorphous silicates in the disk and outflow sources and is also directly connected to the strength of the features. The outflow sources have a lower abundance of crystalline silicates.\\\\ \\\\ \\noindent{\\em 23 micron complex:}\\\\ The strength of the 23.0 micron feature with respect to the strength of the 23.7 micron feature is higher in the outflow sources. The 24.5 micron feature is more prominent in the disk sources. In Paper III we will demonstrate that in the (massive) outflow sources, the ratio of enstatite over forsterite is higher. This would imply that the 23.0 micron feature results indeed from enstatite and that although enstatite has a (weak) feature around 24.5~$\\mu$m, another material is (mainly) responsible for the 24.5 micron feature.\\\\ \\\\ \\noindent{\\em 28 micron complex:}\\\\ The 28 micron complex differs significantly from source to source and several outflow sources are affected by a lack of data. Therefore, no clear difference could be pointed out. \\\\ \\\\ \\noindent{\\em 33 micron complex:}\\\\ The features appear sharper in the outflow sources than in the disk sources. This might have to do with the formation history of the crystalline silicates (see Paper III). The (post-)Red Supergiants separate out by almost equally strong 32.8 and 33.6~micron features, even after subtracting the spurious 32.97~micron feature. This is not a characteristic of all outflow sources.\\\\ \\\\ \\noindent{\\em 40 micron complex:}\\\\ For the 40 micron complex the main difference is found in the strength ratio of the 43.0 and 43.8~micron features. In the disk sources, the 43.8~micron feature is more prominent and in the outflow sources the 43.0~micron feature. This has to do with the high abundance of crystalline H$_2$O-ice in the outflow sources.\\\\ \\\\ \\noindent{\\em 60 micron complex:}\\\\ The main difference in this complex is found in the peak position of the 62~micron feature. This feature is probably a blend of crystalline H$_2$O-ice, peaking around 60~$\\mu$m (the outflow sources), and diopside plus possibly other crystalline silicates, peaking around 65~$\\mu$m (the disk sources). Also, the 69.0~micron feature is much stronger in the disk sources than in the outflow sources, which is probably an abundance effect.\\\\ \\\\ It is likely that the different history of these dust grains are responsible for the spectral differences. MYW suggested that low temperature crystallization has taken place in the disk sources. In contrast, the crystallization of silicates in outflow sources is likely to have taken place close to the star; i.e., high temperature crystallization.  Most features can be explained with forsterite and enstatite. It is clear that several other features lack a proper identification, and also the strength ratios are not always correct, suggesting the presence of more contributing materials. Diopside has been mentioned but the strength of several features attributed to diopside suggests that it is only a minor component.   The growth of enstatite crystals can occur in a preferential direction of one of the crystallographic axes. In chondritic porous interplanetary dust Particles (IDPs), which are believed to be the most pristine material of our solar system, enstatite has been found in which one or 2 directions of the crystallographic axes were severely depressed (Bradley et al, 1983). These enstatite crystals are likely formed from the vapour phase in a low density hot gas environment (Bradley et al, 1983 and references therein), in conditions which are similar to the conditions in the outflows of stars. It is therefore possible that one or two of the crystallographic axes ([010] and [001] according to what has been seen in IDPs) are depressed. This will change the emission spectrum of this material in such a way that certain feature will be weak, corresponding to the depressed crystallographic axis. This scenario might help to understand why the relative strength of some features differ from the lab spectra.  Another abundant species is H$_2$O ice. Kouchi and Kuroda (1990) demonstrate that below 70~K the ice rapidly turns amorphous under the bombardment of UV- photons. Calculations based on their results and reasonable assumptions for the UV flux of the central stars would imply that all the crystalline water ice turns amorphous within a couple of days. Apparently the ice is efficiently shielded from the UV radiation of these stars. Also the low temperature of the ice is an indication for the shielding of the crystalline water-ice. If the material is very clumpy as is seen in e.g. the Helix Nebula (first seen by Baade and reported by Verontsov-Velyaminov (1968), it might provide a suitable environment to survive the UV radiation.  In the outflow sources the abundance of crystalline water ice seems much higher than in the disk sources. Assuming that initially the same amount of H$_2$O was formed, this difference points to a scenario in which crystalline water ice is destroyed on time scales longer than the outflow time scales. In the outflow of stars, density enhancements has been found, where molecules and dust are protected from the harsh UV radiation. This might explain the crystalline water-ice content even for the cold dust. Also inside disks particles are protected against the UV radiation, but the turbulence in disks likely causes that most dust particles will ultimately be exposed to the destructive UV radiation, long before the disk has been blown away.  With about 80\\% of the features identified, see Table~\\ref{tab:complex} for the identifications, the age of astromineralogy has now really started. Still, about 20\\% of the features lack a proper identification. There is also ample evidence that for a significant amount of the `identified' features more dust components are necessary to explain the relative strength of the infrared bands. New laboratory measurements of cosmic relevant dust species up to at least 100~$\\mu$m are required to identify these features. We can now start to fully exploit the historical information of the conditions the dust experienced, by studying the composition and chemical structure of the grains. The differences found between the different laboratory datasets, are a significant obstacle in the quantitative interpretation of the spectra. The properties of the sample (e.g. the stoichiometry, size of the particles, number and size of the individual crystals in a particle) should all be well checked before and after the measurents of the optical constants. We would also recomend that the different astrophysical laboratories exchange their samples in order to have a better understanding of the possible sources of differences in their datasets.   \\vspace{1.0cm} \\noindent{\\bf Acknowledgements.}\\\\ FJM wants to acknowledge the support from NWO under grant 781-71-052 and under the Talent fellowship programm. LBFMW acknowledges financial support from an NWO `Pionier' grant."
    },
    "0201/astro-ph0201132_arXiv.txt": {
        "abstract": "{Recent investigations \\citep[e.g][]{Han-2001:a} have shown that fitting the Hipparcos observations with an orbital model when the astrometric wobble caused by the companion is below the noise level can have rather unexpected consequences.  With new astrometric missions coming out within the next ten years, it is worth investigating the orbit reconstruction capabilities of such instruments at low signal-to-noise ratio.  This is especially important because some of them will have no input catalogue thus meaning that all the orbital parameters will have to be derived from scratch. The puzzling case of almost parabolic orbits is also investigated. ",
        "introduction": "Among the 23882 binaries of the Hipparcos Catalogue \\citep{Hipparcos}, only 235 were processed with an orbital model (with seven additional parameters), the so-called DMSA/O entries.  However, the seven parameters were all fitted for only 45 of them.  In the other cases, the value of some parameters were assumed from previous investigations, mainly spectroscopic and interferometric orbits.  If the same proportion is assumed for GAIA, one will end up with about $5\\,10^5$ systems for which one needs to derive an orbit from scratch.  However, simulations have shown that the number of systems observed by GAIA for which an orbit is worth deriving is about $10^{7}$ \\citep{GAIA-RB}.  The significant improvement of the fit of DMSA/O observations when the orbital model is adopted does not mean they can constrain all the seven parameters.  Assuming the value of some parameters was not just useful, it was sometime necessary in order to derive a realistic orbit.  For instance, HIP~85749 is a well known spectroscopic binary \\citep{Lucke-1982:a} with a period of 418 days.  Though almost three orbital revolutions took place during the Hipparcos mission and an orbital model is well appropriate, a 7-parameter $\\chi^2$ fit leads to $e\\approx 0.99$ whereas the radial velocities yield 0.21.  We report on the robustness of the fit of the seven parameters at different levels of noise on two dimensional and one dimensional (Hipparcos-like) observations.  From now on, it will be assumed that these observations have already been corrected for the parallactic effect and the proper motion, i.e. one will only deal with relative positions whose origin (center of mass or primary) is at rest.  Unlike \\citet{Sozzetti-2001:a}, we limit ourselves to $S/N\\sim 1$. ",
        "conclusions": "\\label{Sect:Conclusions}  One cannot expect to do as much and as well with 1D observations as with 2D ones.  In between these two extremes, the second coordinate with a lower precision is not worth getting unless the ratio of the two precisions does not exceed 5.  The $\\chi^2$ minimization can advantageously be replaced with a 3-stage evaluation thus yielding a global search in a 2-dimensional ($e$, $T$) space.  The period can indeed be independently guessed using period-search technique.  Owing to this low dimension, a grid approach is shown to be quite efficient even if one cannot prevent some false solutions to show up.  The results by \\citet{Han-2001:a} have already warned the community about the reliability of the astrometric orbits, especially the inclinations, at low S/N even when some orbital parameters are adopted from spectroscopy \\citep{Pourbaix-2001:a,Pourbaix-2001:b}.  The situation is likely to get worse when all the orbital parameters are derived from scratch.  All these improvements in the way the orbits are derived will remain useless unless one first finds a criterion for assessing the actual constraint on the projected areal constant and therefore on $i$ and $e$."
    },
    "0201/astro-ph0201242_arXiv.txt": {
        "abstract": "{ We report the results of a detailed analysis of the  temperature structure  of the X-ray emitting plasma halo of M~87, the cD galaxy of the Virgo Cluster. Using  the MEKAL model, the data provide strong indications that the intracluster medium has a single phase structure locally, except the regions associated to the radio structures. The deprojected spectrum at each radius is well fitted by a single temperature MEKAL model, except for the very central region ($<$  2 arcmin) which seems to be affected by the jet and radio lobe structure. The temperature of the intracluster plasma is 1 keV at the center and gradually increases to 2.5 keV at 80 kpc. We have also fitted spectra using the APEC code. Although the large changes of the strength of K$\\alpha$ lines causes a  discrepancy between the Fe-L and Fe-K lines for the APEC results, the overall temperature structure  has not changed. There is no sign of excess absorption in the spectral data. The single-phase nature of the intracluster medium is in conflict with the standard cooling flow model which is based on a multi-phase temperature structure. In addition, the signature of gas cooling below 0.8 keV to zero temperature is not observed as expected for a cooling flow. The gravitational mass profile derived from the temperature and density distribution of the intracluster gas shows two distinct contributions that can be assigned to the gravitational potential of the cD galaxy and the cluster. The central temperature of the intracluster medium agrees well with the potential depth and the velocity dispersion of the cD galaxy. The latter result implies that the central region of the intracluster medium is equivalent to a virialized interstellar medium in M 87. ",
        "introduction": "In the  cores of many clusters of galaxies, X-ray imaging data show a highly peaked surface brightness profile (e.g. Fabian et al. 1981). The radiative cooling time in these regions is much less than a Hubble time. Without a heating process, the gas cools to low temperature and results in a ``cooling flow'' (Fabian 1994 for a review). The mass flow rate, \\mdot, that is deduced in the standard cooling flow model, is approximately proportional to the radius. This implies that matter is deposited throughout the entire cooling flow region. It also implies that the gas in the cooling flow zone is ``multi-phase'' on scales small enough that the inhomogeneities have escaped the observation so far.  ASCA and ROSAT observations confirmed the existence of cooler gas in the cores of cooling flow clusters as expected (e.g. Allen \\& Fabian 1994, Ikebe et al. 1999, Ikebe 2001). The presence of an  intrinsic absorber in excess of the galactic absorption, as it might be expected to arise from the accumulation of cold gas in the cooling flows, was also inferred from the analysis of the ASCA spectra (e.g. Allen 2000; Allen et al. 2001). In fact, the fitting of the ``multi-phase'' cooling flow model to the observed spectra  including temperature phases that cool below the X-ray emitting temperature regime, requires the inclusion of excess absorption to  satisfactorily reproduce the observed spectra. It has also been argued, that inclusion of excess absorption leads to values of the mass flow rate as obtained from imaging data, $\\dot{M}_{\\rm{I}}$  in good agreement with those obtained from spectral analysis, $\\dot{M}_{\\rm{S}}$ (Allen 2000).  However, without  assuming excess absorption, the values of $\\dot{M}_{\\rm{S}}$ are systematically lower than $\\dot{M}_{\\rm{I}}$ (e.g. Ikebe et al. 1999; Makishima et al. 2001). Instead of the cooling flow model, a two temperature spectral model can well fit ASCA spectra of the cooling flow clusters. The temperatures of the softer component are about half of the hotter component (Ikebe 2001). The central temperatures of some cD galaxies obtained with the ROSAT PSPC are close to those obtained from normal elliptical galaxies with the same stellar velocity dispersion (Matsushita 2001). Since cooling flow clusters always possess central dominant galaxies, the cool components may reflect the potential structure, rather than cooling gas (Ikebe et al. 1999, Makishima et al. 2001, Ikebe 2001).  Recently, it was discovered with the RGS instrument onboard XMM-Newton that there is little X-ray emission from a component with a temperature below a certain lower cutoff value that differs from object to object (e.g. Tamura et al. 2001; Kaastra et al. 2001). However, with the RGS one observes only the very central part of the cluster.  M~87 is the cD galaxy of the nearest rich cluster of galaxies, the Virgo Cluster. It is very luminous in X-rays, and  is suggested to have a ``cooling flow'' with \\mdot~of about  10 $M_\\odot\\rm{yr^{-1}}$ (Stewart et al. 1984;  Fabian et al. 1984). M~87 also hosts a central active galactic nuclei (AGN). The radio emission is complex and there are two strong  lobe structures (e.g. B\\\"ohringer et al. 1995). An enhancement of the X-ray emission around the lobes  was discovered with the EINSTEIN HRI (Feigelson et al. 1987) and ROSAT PSPC (B\\\"ohringer et al. 1995)  In  addition to the XMM RGS results mentioned above, the XMM-Newton observatory offers the possibility to perform a very detailed spectral and spatial analysis of the cooling flow area. This is in particular important since the RGS spectrum taken at the very central part of the intracluster medium (ICM) may be strongly affected by resonance line scattering and thus the interpretation of the results requires the combined analysis of spectra taken at different radii from the center. Early results of the XMM observation of M 87 have  already been published (B\\\"ohringer et al. 2001; Belsole et al. 2001). The projected spectrum is fitted with a single temperature model (B\\\"ohringer et al. 2001). The abundance drop at the central region is interpreted as resonance scattering. The detailed spectral analysis on the jets and  the radio lobe structure is shown in Belsole et al. (2001).  In this paper, we report on the more detailed temperature structure using a deprojecting analysis, and apply a cooling flow model to the spectral analysis. We employ a distance of M~87 of 17 Mpc. We adopt for the solar iron abundance the `photospheric' value, Fe/H $=4.68\\times 10^{-5}$ by number (Anders \\& Grevesse 1989). Unless otherwise specified,  we use 90\\% confidence error regions. ",
        "conclusions": "The major result of this paper is strong evidence that the ICM in the halo of M 87 is isothermal locally, except probably for the very central region, where more than one temperature component is present. This could be partly due to the interaction effects of the ICM with the jet and the radio lobes of M 87. This finding and the fact that no spectral signature of low temperature components below temperatures of 0.8 keV are observed is in disagreement with the standard cooling flow model which predicts a strongly multi-phase structure of the ICM and a distinct temperature distribution and resulting spectrum with clear signature of low temperature components. We have also shown that strong inhomogeneities in the metal distribution cannot resolve this problem. Therefore we conclude that the scenario for the dense gas with short cooling time in the centers of clusters needs to be revised. Probably heating processes that can substantially reduced the mass deposition have to be reconsidered, like the energy input of the central AGN. The central temperature closely reflects the gravitational potential depth of the central galaxy, rather than the  existence of a cooling flow."
    },
    "0201/astro-ph0201060_arXiv.txt": {
        "abstract": "One of the major challenges to identification of the 3.3, 6.2, 7.7, 8.6, and 11.3$\\micron$ interstellar infrared (IR) emission bands with polycyclic aromatic hydrocarbon (PAH) molecules has been the recent detection of these bands in regions with little ultraviolet (UV) illumination, since small, neutral PAH molecules have little or no absorption at visible wavelengths and therefore require UV photons for excitation.  We show here that our ``astronomical'' PAH model, incorporating the experimental result that the visual absorption edge shifts to longer wavelength upon ionization and/or as the PAH size increases, can closely reproduce the observed infrared emission bands of vdB 133, a UV-poor reflection nebula.  It is also shown that single-photon heating of ``astronomical'' PAHs in reflection nebulae near stars as cool as $\\Teff=3000\\K$ can result in observable emission at 6.2, 7.7, 8.6, and 11.3$\\micron$. Illustrative mid-IR emission spectra are also calculated for reflection nebulae illuminated by cool stars with $\\Teff=3500, 4500, 5000\\K$. These will allow comparison with future Space Infrared Telescope Facility (SIRTF) observations of vdB 135 ($\\Teff=3600\\K$), vdB 47 ($\\Teff=4500\\K$), and vdB 101 ($\\Teff=5000\\K$).  It is also shown that the dependence of the 12$\\micron$ IRAS emission relative to the total far-IR emission on the effective temperature of the exciting star is consistent with the PAH model expectation for $3000\\K \\le \\Teff \\le 30000\\K$. ",
        "introduction": "} Since their first detection in the planetary nebulae NGC 7027 and BD+30$^{\\rm o}$3639 (Gillett, Forrest, \\& Merrill 1973), the distinctive set of infrared (IR) emission features at 3.3, 6.2, 7.7, 8.6, and 11.3$\\micron$ have been observed in a variety of objects with a wide range of physical conditions, including planetary nebulae, protoplanetary nebulae, reflection nebulae, HII and ultracompact HII regions, circumstellar envelopes, the diffuse interstellar medium (ISM) of the Milky Way Galaxy, and external galaxies (see Tielens et al.\\ 1999 and Sellgren 2001 for recent reviews). Although the exact nature of their carriers remains unidentified -- they remain known as ``the Unidentified Infrared (UIR) bands'' -- it is now widely thought that they originate from some sorts of aromatic hydrocarbons. Various carbonaceous materials have been proposed as the UIR band carriers. In general, they can be categorized into two classes: (1) pure, free-flying aromatic molecules -- polycyclic aromatic hydrocarbon molecules (PAHs; L\\'{e}ger \\& Puget 1984; Allamandola, Tielens, \\& Barker 1985); (2) carbonaceous grains with a partly aromatic character -- hydrogenated amorphous carbon (HAC; Duley \\& Williams 1981; Jones, Duley, \\& Williams 1990), quenched carbonaceous composite (QCC; Sakata et al.\\ 1990), coal (Papoular et al.\\ 1993), fullerenes (Webster 1993), and surface-graphitized nanodiamonds (Jones \\& d'Hendecourt 2000).  In most current models, the UIR band emission involves three sequential steps: (1) excitation by absorption of an energetic starlight photon (usually ultraviolet [UV]); (2) rapid ($\\sim 10^{-12}-10^{-10}\\s$) redistribution of all or part of the absorbed photon energy over all available vibrational modes; (3) radiative relaxation via IR fluorescence.  Among the existing proposed carriers, the PAH model is gaining increasing acceptance because of (1) the close resemblance of the UIR spectra (frequencies and relative intensities) to the vibrational spectra of PAH molecules (e.g. see Allamandola, Hudgins, \\& Sandford 1999) and (2) the ability of a PAH molecule to emit efficiently in the UIR wavelength range following single photon heating (Greenberg 1968; L\\'{e}ger \\& Puget 1984; Allamandola, Tielens, \\& Barker 1985; Draine \\& Li 2001). In contrast, larger carbonaceous grains are unlikely to be efficient UIR emitters since the timescale for the absorbed photon energy to diffuse ($\\sim 10^{-9}\\s$) is much shorter than the IR vibrational emission timescale ($\\sim 0.1\\s$; see Tielens et al.\\ 1999).  Early observations of the UIR emission bands were made in regions with strong UV irradiation, and PAH excitation was expected since all PAH species are strongly absorbing in the vacuum ultraviolet ($\\lambda \\simlt 3000$\\AA). The PAH UV excitation model has recently been challenged by the ISO ({\\it Infrared Space Observatory}) detection of the UIR bands in $\\vdb$, a reflection nebula illuminated by a binary system with little UV radiation (see Uchida, Sellgren, \\& Werner 1998). The UIR spectrum of this UV-poor region closely resembles those observed in sources with much harsher UV environments (Uchida et al.\\ 2000). This appears to be in conflict with the view that PAHs are primarily excited by UV photons, as would be expected based on laboratory studies showing that the absorption by {\\it small, neutral} PAHs has a sharp cut-off in the UV, with little or no absorption in the visible (see Sellgren 2001 for a review). Uchida et al.\\ (2000) obtained 5--15$\\micron$ ISOCAM spectra of 6 reflection nebulae and found no systematic spectroscopic differences despite values of $\\Teff$ ranging from 6800$\\K$ to 19000$\\K$.  The PAH electronic absorption edge is known to shift to longer wavelength with increasing size and/or upon ionization (see Allamandola et al.\\ 1989, Salama et al.\\ 1996 and references therein). While the largest experimentally studied PAH molecule to date is dicoronylene C$_{48}$H$_{20}$ (Allamandola, Hudgins, \\& Sandford 1999), astronomical PAHs are believed to be larger (e.g., the mean size for the Milky Way PAHs is $\\approx 6$\\AA, corresponding to $N_{\\rm C}\\approx 100$ [Li \\& Draine 2001a]). The astronomical PAH model -- with the size/ionization dependence of the PAH absorption edge taken into account (see \\S A2 in Li \\& Draine 2001a) -- is successful in explaining the observed mid-IR spectra of the Milky Way diffuse ISM (Li \\& Draine 2001a), the quiescent molecular cloud SMC B1\\#1 in the Small Magellanic Cloud (Li \\& Draine 2001b), and the UIR band ratios for a wide range of environments ranging from reflection nebulae, HII regions, photodissociation regions (PDRs), molecular clouds in the Milky Galaxy to normal galaxies, starburst galaxies, and a Seyfert 2 galaxy (Draine \\& Li 2001).  In this {\\it Letter} we show that the astronomical PAH model is consistent with the observed UIR emission from UV-poor environments. In \\S\\ref{sec:vdb133} we verify that the astronomical PAH model can quantitatively reproduce the vdB 133 UIR spectrum. We further demonstrate in \\S\\ref{sec:coolrn} that the UIR bands are also expected for reflection nebulae which are even more UV-poor than vdB 133, and we provide model spectra for comparison with future {\\it Space Infrared Telescope Facility} (SIRTF) observations of reflection nebulae near cool stars. In \\S\\ref{sec:12um} we show that the predicted $\\Teff$ dependence of the ratio of the IRAS 12$\\micron$ emission to the total far-IR emission is also consistent with observations. Our results are discussed in \\S\\ref{sec:discuss}, and our conclusions are summarized in \\S\\ref{sec:sum}. ",
        "conclusions": "}  We have modeled the excitation of PAH molecules in UV-poor regions. It is shown that the astronomical PAH model provides a satisfactory fit to the UIR spectrum of vdB 133, a reflection nebulae with the lowest ratio of UV to total radiation among reflection nebulae with UIR bands detected. It is also shown that astronomical PAHs can be pumped by cool stars with even less UV radiation. It is further shown that the PAH model predicts a dependence of $\\irratio$ for reflection nebulae which is consistent with observations for $3000\\K \\le \\Teff \\le 30000\\K$. We conclude that PAHs appear able to account for the UIR band emission observed in reflection nebulae."
    },
    "0201/astro-ph0201256_arXiv.txt": {
        "abstract": "We analyze PLANET and MACHO observations of \\mb, the first nearly-normal microlensing event for which high signal-to-noise-ratio data reveal a well-covered, short-duration anomaly. This anomaly occurs near the peak of the event. Short-duration anomalies near the peak of otherwise normal events are expected to arise both from extreme-separation (either very close or very wide), roughly equal-mass binary lenses, and from planetary systems. We show that the lens of \\mb\\ is in fact an extreme-separation binary, not a planetary system, thus demonstrating for the first time that these two important classes of events can be distinguished in practice. However, we find that the wide-binary and close-binary lens solutions fit the data equally well, and cannot be distinguished even at $\\Delta\\chi^2=1$. This degeneracy is qualitatively much more severe than the one identified for \\smc\\ because the present degeneracy spans two rather than one dimension in the magnification field and does not require significantly different blending fractions. In the appendix, we explore this result, and show that it is related to the symmetry in the lens equation. ",
        "introduction": "The hallmark of planetary microlensing events is a short deviation from an otherwise normal, point-source/point-lens (hereafter PSPL) event.  \\citet{mp} showed that extrasolar planets could be detected from such events, and \\citet{gl} gave an explicit prescription for how the planet/star mass ratio $q$ ($\\ll 1$) and the angular separation $d$ (in units of the angular Einstein radius $\\theta_{\\rm E}$) could be reconstructed by decomposing the event light curve into its ``normal'' and ``perturbed'' components.  Work during the ensuing decade has elucidated many additional subtleties of planetary light curves, but their fundamental characterization as briefly perturbed PSPL events has remained intact.   Of particular importance in the present context, \\citet{gs} showed that events with small impact parameter ($u_0\\ll 1$; where $u_0$ is the minimum separation between the source and the lens center of mass in units of $\\theta_{\\rm E}$) probe the so-called ``central caustic'' of the lens geometry, making them much more sensitive to the presence of planets than the larger impact-parameter events analyzed by \\citet{gl}, which probe the outer ``planetary caustic''.  These central-caustic events are of exceptional importance, even though they are intrinsically rare.  They are rare simply because the central-caustic is much smaller than the planetary caustic, so the great majority of planet-induced deviations (of fixed fractional amplitude) are due to planetary caustics.  However, the probability of detecting a planet is much greater in small impact-parameter events, partly because the source is guaranteed to pass close to the central caustic, which almost by definition is near the center of the lens geometry ($\\vect{u}=0$) and partly because even the sensitivity of planetary caustics is enhanced for $u_0\\ll 1$.  (Here $\\vect{u}$ denotes the source position on the sky, normalized to $\\theta_{\\rm E}$, with respect to the lens center of mass.)  By contrast, higher impact-parameter events miss the central caustic, and they are likely to miss the planetary caustic as well because it lies in a random position relative to the source trajectory.  Because of their higher sensitivity to planets, and because they can be recognized in real time,  low impact-parameter events are monitored more intensively than typical events by microlensing follow-up networks, which in turn further enhances their sensitivity to planets.  Central-caustic events, like their planetary-caustic cousins, involve a short deviation from an otherwise normal PSPL light curve. The major difference between these two classes of planetary events is that central-caustic anomalies always occur near the peak, whereas planetary-caustic perturbations can occur anywhere on the light curve, and are typically expected on the wings of the light curve. Of particular importance, for central-caustic events, there is no simple prescription for extracting $d$ and $q$ by decomposing the light curve into ``normal'' and ``perturbed'' components and it is unclear to what degree these two parameters are degenerate.  Another important, albeit accidental, discovery was that planets could give rise to perturbations that are not short compared to the event timescale.  In the course of their search for planetary perturbations among 43 approximately PSPL events, \\citet[see also \\citealt{letter}]{planetsearch} found one event, \\obp, that was asymmetric in a way that was consistent with the presence of a planet.  They argued, however, that the asymmetry was also consistent with parallax effects induced by the Earth's motion around the Sun of the type analyzed by \\citet*{gmb}, and that in general it would be extremely difficult to distinguish between the two possible causes of such an asymmetry. They concluded that, in most cases, microlensing searches are not able to distinguish between parallax and a weak, asymmetric planetary perturbation, and consequently, all such ``detections'' should be ignored. This reduces the sensitivity of microlensing searches to planets, but only by an extremely small amount since, as \\citet{planetsearch} showed, long-timescale asymmetric perturbations account for less than $\\sim 1\\%$ of all planetary events. Hence, the long-timescale asymmetric events also confirm in a way the basic paradigm: planetary perturbations have short durations relative to the parent light curve, and in the rare cases for which they do not, they are not recognizable as a planetary perturbation anyway.  However, not all short timescale deviations are due to planets and therefore the mere detection of such an anomaly does not prove the presence of a planet. \\citet{gg} showed that close binaries ($d\\ll 1$) give rise to light curves that are virtually identical to PSPL events, except when the source comes very close to the lens center of mass ($|\\vect{u}|\\ll 1$). Hence, for events with $u_0\\ll 1$, the light curve looks ``normal'' except for a brief deviation near the peak.  Qualitatively, this is exactly the same behavior as that of central-caustic planetary events. Similarly, light curves of wide-binary ($d\\gg 1$) events can also take the same form if one -- and only one -- of the caustics lies very close to the source's passage. Indeed, a close correspondence between a certain pair of close-binary and wide-binary events was discovered both theoretically \\citep{d99} and observationally \\citep{comp,smc}. It remains an open question under what conditions these various types of events can be distinguished from one another.  If central-caustic planetary events could not be distinguished from close- and/or wide-binary events, it would seriously undermine planet searches in high-magnification events and hence would call into question the basic strategy adopted by microlensing follow-up groups \\citep[e.g.,][]{PLANET}.  Here we analyze the light curve of the microlensing event \\mb, the first intensively monitored event with a short-lived, high signal-to-noise-ratio deviation from an otherwise normal PSPL light curve. We identify the lens as an extreme-separation binary rather than a planetary system, thereby showing that, at least in this case, the two classes can be clearly distinguished. We also show that both wide- and close-binary solutions fit the data equally well, implying that although planetary perturbations can be distinguished from those arising from extreme-separation binaries, the discrimination between very close and very wide binaries may be difficult in practice. ",
        "conclusions": "\\mb\\ is the first microlensing event with a short-lived, high signal-to-noise-ratio anomaly, characteristics that could betray the existence of a planet around the lensing star. Nevertheless, we conclude that the lens of \\mb\\ is not a planetary system, but an extreme-separation (very close or very wide) binary composed of components of similar mass, based on the result of the light curve fit as well as the extreme value of event duration and blending fraction required for any plausible ``planetary'' fit."
    },
    "0201/astro-ph0201126_arXiv.txt": {
        "abstract": "We study the transition from npe-type nuclear matter (consisting of neutrons, protons, and electrons) to matter containing strangeness, using a Walecka-type model predicting a first-order kaon-condensate phase transition. We examine the free energy of droplets of K-matter as the density, temperature, and neutrino fraction are varied.  Langer nucleation rate theory is then used to approximate the rate at which critical droplets of the new phase are produced by thermal fluctuations, thus giving an estimate of the time required for the new (mixed) phase to appear at various densities and various times in the cooling history of the proto-neutron star.  We also discuss the famous difficulty of ``simultaneous weak interactions'' which we connect to the literature on non-topological solitons.  Finally, we discuss the implications of our results to several phenomenological issues involving neutron star phase transitions. ",
        "introduction": " ",
        "conclusions": "\\label{sec5}  We have, then, an intriguing possible scenario in which the PNS manages to settle into its ground state without forming the kaon-condensate mixed phase which is the true ground state of the system.  At high initial temperatures, the nucleation of K-matter droplets is suppressed by the presence of neutrinos, even though at these high temperatures the seeding of droplets is in principle fast.  Over a wide range of densities kaon droplets are produced copiously by thermal fluctuations, but they are not yet stable due to the increase in effective critical density.  As the star cools, the restriction coming from the presence of neutrinos is relaxed, but the intrinsic fluctuation rate drops.  For an initial PNS core density that is not too far above the nominal ($T=0$) critical density for the formation of a K-matter mixed phase, therefore, it is likely that the star will cool not into the true ground state, but, rather, into a meta-stable state consisting of electrically neutral npe matter.  This scenario may be relevant to understanding various phenomenological issues.  For example, the apparent existence of anomalously heavy neutron stars with masses $M \\sim 2 M_{sun}$ \\cite{zhang,kerk,orosz} might be explained by the anomalously stiff equation of state of npe-type matter relative to matter with a kaon condensate.  Generally, if the various possible phase transitions thought to occur in neutron star matter can be avoided by the impossibility (or, equivalently, extreme slowness) of nucleation of the new phase, a relatively stiff equation of state may sometimes be maintained over a more extended range of density than would be expected naively. Hence, the existence of such heavy stars may not be sufficient evidence to rule out the existence of kaon-condensation at $3-5 n_0$, especially if these stars were born with smaller masses and only subsequently (that is, once cold) acquired larger masses via accretion from companion stars.  Additionally, metastability of the sort introduced above may potentially be useful in understanding the properties of GRB's or other poorly-understood explosive events.  A relatively light PNS, as we have argued, may cool and deleptonize without the K-matter mixed phase forming, even when the star's core density exceeds the critical density for the transition.  As is evident from Fig. \\ref{fig6}, however, the thermal kaon density increases monotonically and steeply with density, so that, even at the very low temperatures $T \\ll 1$ MeV eventually attained in the neutron star, there is some density at which seeding may become possible in a reasonable time. At very least, with increasing density, one eventually encounters the second-order point at which the kaon field may be produced smoothly with no need for seeding.  Hence, if an initially metastable neutron star begins to accrete matter from a companion binary (or, additionally, if an initially rotating metastable neutron star gets spun down via accretion) the central density may increase sufficiently for K-matter to begin to appear.  Once the kaon matter appears, however, there will be a feedback effect, due to the softening of the equation of state.  The production of a small quantity of K-matter in the core would allow the star to contract slightly, thus increasing the density in the core, and increasing the size of the region in which kaonic matter can appear.  Further kaon production leads to further collapse, and vice versa.  Thus if mass accretion and/or spin down results in the critical density for the onset of a second-order kaon-condensate transition being reached in the neutron star core (or, equivalently, if one reaches the low-temperature effective critical density at which a mixed phase can be nucleated spontaneously with sufficient speed) one would expect an explosive event in which the star contracts significantly, resulting in the release of a tremendous amount of energy (of order tenths of $M_{sun}$). Cheng and Dai have discussed a similar proposal in which accretion-induced conversion to strange quark matter is suggested as a possible explanation for Gamma-Ray Bursters. \\cite{chengdai}  One especially interesting aspect of such a collapse is its potentially turbulent nature.  The picture is of a pure K-matter core seeding the mixed phase through several kilometers of material above it.  In effect, the K-matter boils off of the outer edge of the second-order core and floats upward to form the mixed phase throughout the entire region of the mixed phase's energetic favorability.  This implies an upward and downward transfer of matter that closely resembles turbulent convection, but in which strangeness rather than heat is the substance being convected.  As mentioned at the end of the previous section, more reliable calculations need to be performed in order to better understand how the slowness of the weak interactions affect the original nucleation rate estimates based on Langer's formula.  We have argued that the correct framework for these future calculations involves the formalism of quantum tunneling (or thermal activation) in a theory with a (nearly) conserved global charge representing strangeness.  It is also worth mentioning that the scenario outlined here is the most realistic potential application of the Q-ball nucleation formalism developed in \\cite{lee1,lee2}; theorists up to now have relied on supersymmetric models to find possible theories containing charged scalars supporting non-topological solitons.  \\begin{center}  \\large{ \\textbf{Acknowledgements}}  \\end{center} Sanjay Reddy, Eduardo Fraga, and Guy Moore are acknowledged for helpful discussions, though any inaccuracies in the present paper are purely the responsibility of the author.  This work is supported in part by a National Science Foundation Graduate Research Fellowship and by the US Department of Energy grant DE-FG03-00ER41132.   \\newcommand{\\IJMPA}[3]{{ Int.~J.~Mod.~Phys.} {\\bf A#1}, (#2) #3} \\newcommand{\\JPG}[3]{{ J.~Phys. G} {\\bf {#1}}, (#2) #3} \\newcommand{\\AP}[3]{{ Ann.~Phys. (NY)} {\\bf {#1}}, (#2) #3} \\newcommand{\\NPA}[3]{{ Nucl.~Phys.} {\\bf A{#1}}, (#2) #3 } \\newcommand{\\NPB}[3]{{ Nucl.~Phys.} {\\bf B{#1}}, (#2)  #3 } \\newcommand{\\PLB}[3]{{ Phys.~Lett.} {\\bf {#1}B}, (#2) #3 } \\newcommand{\\PRv}[3]{{ Phys.~Rev.} {\\bf {#1}}, (#2) #3} \\newcommand{\\PRC}[3]{{ Phys.~Rev. C} {\\bf {#1}}, (#2) #3} \\newcommand{\\PRD}[3]{{ Phys.~Rev. D} {\\bf {#1}}, (#2) #3} \\newcommand{\\PRL}[3]{{ Phys.~Rev.~Lett.} {\\bf {#1}}, (#2) #3} \\newcommand{\\PR}[3]{{ Phys.~Rep.} {\\bf {#1}}, (#2) #3} \\newcommand{\\ZPC}[3]{{ Z.~Phys. C} {\\bf {#1}}, (#2) #3} \\newcommand{\\ZPA}[3]{{ Z.~Phys. A} {\\bf {#1}}, (#2) #3} \\newcommand{\\JCP}[3]{{ J.~Comput.~Phys.} {\\bf {#1}}, (#2) #3} \\newcommand{\\HIP}[3]{{ Heavy Ion Physics} {\\bf {#1}}, (#2) #3} \\newcommand{\\RMP}[3]{{ Rev. Mod. Phys.} {\\bf {#1}}, (#2) #3} \\newcommand{\\APJ}[3]{{Astrophys. Jl.} {\\bf {#1}}, (#2) #3}"
    },
    "0201/astro-ph0201310_arXiv.txt": {
        "abstract": "Synchrotron emission of relativistic particles in magnetic fields is a process of paramount importance in astrophysics. Although known for over thirty years, there are still aspects of this radiative process that have received little attention, mainly because they appear only in extreme conditions. In the present paper, we first provide a general introduction to synchrotron emission, using a formalism that represents a generalization of the standard calculations. The use of this formalism allows us to discuss situations in which charged particles can radiate coherently, with special attention for the cases in which the production occurs in the form of a bunch of particles created in a pulse of very short duration. We calculate the spectra of the radiation for both monoenergetic particles and distributions of particles with different Lorentz factors. For both cases we study the conditions for the coherent effects to appear, and demonstrate that in the limit of incoherent emission we reobtain the well known results. ",
        "introduction": "Synchrotron radiation and its importance for astrophysics have been discussed in such a large number of papers that it is hard to believe there is anything else left to say. The basic reviews are those in Refs. \\cite{GS1,GS2} while a detailed description of the standard theory is presented in \\cite{RL}. Nevertheless, most previous calculations are restricted to conditions that were considered {\\it reasonable} for astrophysical standards. These {\\it reasonable} standards are now considerably different from those of three decades ago, when synchrotron emission was first studied in astrophysics. We now know that there are situations in which standard calculations of synchrotron emission are not applicable. Two examples can be easily found and will be discussed: coherence effects from pulsed bunches of particles and synchrotron backreaction. Although some pieces of work have previously appeared in the literature, in our opinion a complete treatment of these phenomena is still missing. This paper will be devoted to the study of coherent synchrotron emission, in a very general framework, so that the conclusions may be applied to the cases of interest. In an accompanying paper \\cite{paperII} (hereafter paper II) we will discuss the synchrotron backreaction, another topic that is rarely discussed in the literature, and for which a comprehensive treatment is still lacking. In paper II we will adopt the formalism introduced here.  Coherence effects occur when there are well defined phase relations among the radiating particles, so that both intensity and spectra of the resulting radiation suffer from non-negligible interference effects. In these cases, a system of $Z$ particles with Lorentz factor $\\gamma$ has a synchrotron radiation which is up to $Z^2$ times the spectrum of a single particle, to be compared with the incoherent radiation, in which case the emission rate is $Z$ times the emission rate of a single particle.  This is not a new point: there are many papers in which this enhancement of the radiation was pointed out (e.g. \\cite{general,saggion,benford,kirk}). Nevertheless we think that there are important differences between these papers and the present calculations. First, all previous papers that we are aware of discuss the specific case of curvature radiation, mainly because the application kept in mind is that to pulsar radio emission, where it seems that coherence effects may be needed. Second, these past calculations point to the evaluation of the power of the emitted radiation; we will devote part of this paper to point out that in case of impulsive coherent emission, this may be a not well defined quantity. Third, the previous calculations take care of the coherent emission from bunches of particles all with the same Lorentz factor, while in the present paper we generalize the results to the case of a spectrum of radiating particles. As a special case, we recover the previous results.  Strictly speaking the literature that we are aware of deals with the process of curvature radiation, thought to be at work in pulsar magnetospheres. This case is formally similar to the one considered here but physically the conditions for the occurrence of coherent effects may be quite different.  In addition to these points, we propose a new kind of formal approach to the calculations of the synchrotron emission from ensemble of particles. The new approach reproduces the results of the standard approach but also provides new insights on the physical interpretation of those results.  For the cases where coherence effects are expected, we discuss the factors that may be responsible for the decoherence of the emitted radiation, or, in other words, the factors that can transform the emitted radiation from coherent to incoherent.  The paper is structured as follows: in section 2 we describe our formalism for the calculation of synchrotron emission, from which the occurrence of coherent effects arises naturally. In section 3 we describe the concept of radiated power in the case of coherent emission. In section 4 we use the approach introduced in section 2 in order to describe several features of the coherent synchrotron emission from bunches of particles. We also define the condition for coherence to appear. We conclude in section 5. ",
        "conclusions": "We studied the theory of coherent synchrotron emission in the perspective of possible astrophysical applications. The theory is obtained from the general treatment of radiation from an ensemble of particles, so that the usual results are easily recovered when the coherence effects are not relevant.  Our conclusions can be summarized as follows:  {\\it i)} An ensemble of $Z$ monoergetic particles in perfect phase radiate coherently at any frequency, and the total radiated energy is $Z^2$ larger than the energy radiated by a single particle at the same frequency. This can be easily interpreted by recalling that the synchrotron spectra are proportional to the square of the electric charge of a particle: coherent emission from $Z$ particles can be thought as the synchrotron emission of a single particle with charge $Z$.  {\\it ii)} If $Z$ monoenergetic particles have a phase spread $\\Delta\\alpha$, the coherence is limited to frequencies $n\\ll 1/\\Delta\\alpha\\gamma$. At higher frequencies the incoherent result is gradually recovered.  {\\it iii)} In the case of $Z$ particles with a spectrum $N(\\gamma)\\propto \\gamma^{-p}$ in the range $\\gamma_{min}\\leq\\gamma\\leq\\gamma_{max}$ and spread in phase over an angle parametrized as $\\Delta\\alpha=\\xi/ \\gamma_{max}^3$, the condition for coherence is that $n\\ll \\gamma_{max}^2/\\xi$. At high frequencies, $n\\gg \\gamma_{max}^2/\\xi$, the incoherent result is recovered.  {\\it iv)} The coherent emission from a bunch of particles is not likely to be stable in time: in the case of a monoenergetic bunch, small inhomogeneities in the magnetic field structure introduce phase shifts among the particles, so that the coherence condition may be easily broken. In the case of a bunch of particles with a distribution of Lorentz factors, inhomogeneity of the magnetic field and the fact that the Larmor radii of the particles are different makes the stability of the coherence very problematic. In both cases the coherent emission can however be generated at the time of formation of the bunch, for instance in pulsed events.  Despite the difficulty in maintaining the coherent character of the radiation, there seem to be several situations in which invoking the coherence appears to provide the most reasonable explanation for the observations. One case is that of the radio emission from pulsars and is discussed in \\cite{michel}. In this case the radiation is however most likely curvature radiation rather than synchrotron emission. Although the two mechanisms are very similar, there are also technical differences between the two. The second case of possible coherent emission is related to the radio brightness of jets in active galactic nuclei, when the brightness temperature is $T_B>3\\times 10^{17}$ K. This case has been mentioned in the literature (e.g. \\cite{spada,ghisellini}) but never treated in detail.  {\\bf aknowledgments} We are very grateful to F. Pacini and A. Olinto for several constructive discussions and to M. Salvati and T. Stanev for a critical reading of the manuscript. We are also grateful to the anonymous referee for the interesting remarks on the paper."
    },
    "0201/astro-ph0201476_arXiv.txt": {
        "abstract": "Ultraluminous galaxies in the local universe (z$\\leq$0.2) emit the bulk of their energy in the mid and far-infrared. The multiwavelength approach to these objects has shown that they are advanced mergers of gas-rich spiral galaxies. Galaxy-galaxy collisions took place on all cosmological time-scales, and nearby mergers serve as local analogs to gain insight into the physical processes that lead to the formation and trans-formation of galaxies in the more distant universe. Here I review multiwavelength observations --with particular emphasis on recent results obtained with ISO-- of mergers of massive galaxies driving the formation of: 1) luminous infrared galaxies, 2) elliptical galaxy cores, 3) luminous dust-enshrouded extranuclear starbursts, 4) symbiotic galaxies that host AGNs, and 5) tidal dwarf galaxies. The most important implication for studies on the formation of galaxies at early cosmological timescales is that the distant analogs to the local ultraluminous infrared  galaxies are invisible in the ultraviolet and optical wavelength rest-frames and should be detected as sub-millimeter sources with no optical counterparts. ",
        "introduction": "One of the most important discoveries from extragalactic observations at mid- and far-infrared wavelengths has been the identification of a class of ``Luminous Infrared Galaxies'' (LIGs), objects that emit more energy in the infrared ($\\sim${\\ts}5--500{\\ts}$\\mu m$) than at all other wavelengths combined (see \\cite{SandersMirabel} for a comprhensive review).  The first all-sky survey at far-infrared wavelengths carried out in 1983 by the {\\it Infrared Astronomical Satellite} ({\\it IRAS}) resulted in the detection of tens of thousands of galaxies, the vast majority of which were too faint to have been included in previous optical catalogs.  It is now clear that part of the reason for the large number of detections is the fact that the majority of the most luminous galaxies in the Universe are extremely dusty.  Previous assumptions, based primarily on optical observations, about the relative distributions of different types of luminous galaxies---e.g.  starbursts, Seyferts, and quasi-stellar objects (QSOs)---need to be revised.  Galaxies bolometrically more luminous than $\\sim${\\ts}4{\\ts}${\\it L}^*$ (i.e. {\\it L}$_{bol} \\geq 10^{11}{\\ts}{\\it L}_\\odot$) appear to be heavily obscured by dust.  Although luminous infrared galaxies (hearafter LIGs: {\\it L}$_{ir} >{\\ts}10^{11}{\\ts}{\\it L}_\\odot$) are relatively rare objects, reasonable assumptions about the lifetime of the infrared phase suggest that a substantial fraction of all galaxies with {\\it L}$_B$ $>{\\ts}10^{10}${\\it}{\\it L}$_\\odot$ pass through such a stage of intense infrared emission \\cite{Soifer1}.  \\begin{figure}[htb] \\begin{center} \\includegraphics[width=.8\\textwidth]{araa_fig1.eps} \\end{center} \\caption[]{Galaxy luminosity function of Infrared Galaxies compared with other extragalactic objects in the local universe. Among the most luminous galaxies (${\\it L}_{bol} > 10^{11.5}${\\ts}{\\it}L$_\\odot$), infrared galaxies selected from the IRAS survey outnumber optically selected Seyferts and quasars.  For references see \\cite{SandersMirabel}.} \\label{fig1} \\end{figure}  A comparison of the luminosity function of infrared bright galaxies with other classes of extragalactic objects in the local universe is shown in Figure \\ref{fig1}. At luminosities below $10^{11}${\\ts}{\\it L}$_\\odot$, {\\it IRAS} observations confirm that the majority of optically selected objects are relatively weak far-infrared emitters. Surveys of Markarian galaxies confirm that both Markarian starbursts and Seyferts have properties (e.g. ${\\it f}_{60}/{\\it f}_{100}$ and ${\\it L}_{ir} / {\\it L}_B$\\ ratios) closer to infrared selected samples as does the subclass of optically selected interacting galaxies. However because the most luminous galaxies are enshrouded in dust, relatively few objects in optically selected samples are found with ${\\it L}_{ir} > 10^{11.5}${\\ts}{\\it}L$_\\odot$.   \\begin{figure}[htb] \\begin{center} \\includegraphics[width=.8\\textwidth]{araa_fig2.eps} \\end{center} \\caption[]{ Variation of the mean Spectral Energy Distribution (from submillimeter to UV wavelengths) with increasing { L}$_{ir}$ for a 60{\\ts}$\\mu$m sample of infrared galaxies.  ({ Insert}) Examples of the subset ($\\sim${\\ts}15\\%) of ULIGs with ``warm'' infrared color ({ f}$_{25}$/{ f}$_{60}$ $>${\\ts}0.3).  Three objects (1---the powerful Wolf-Rayet galaxy IRAS{\\ts}01002--2238, 2---the ``infrared QSO'' IRAS{\\ts}07598+6508, 3---the optically selected QSO I{\\ts}Zw{\\ts}1) are shown in the inset. For references see \\cite{SandersMirabel}.} \\label{fig2} \\end{figure}   The high luminosity tail of the infrared galaxy luminosity function is clearly in excess of what is expected from the Schechter function. For ${\\it L}_{bol} = 10^{11}-10^{12}${\\ts}{\\it L}$_\\odot$, LIGs are as numerous as Markarian Seyferts and $\\sim${\\ts}3 times more numerous than Markarian starbursts.  Ultraluminous infrared galaxies (hereafter ULIGs: ${\\it L}_{ir} > 10^{12}${\\ts}{\\it L}$_\\odot$) appear to be $\\sim${\\ts}2 times more numerous than optically selected QSOs, the only other previously known population of objects with comparable bolometric luminosities.   \\begin{figure}[htb] \\begin{center} \\includegraphics[width=.8\\textwidth]{araa_fig8_2.eps} \\end{center} \\caption[]{ Well-studied mergers: {\\it (a)}\\ NGC{\\ts}4038/39 (Arp{\\ts}244 = ``The Antennae''); {\\it (b)}\\ NGC{\\ts}7252 (Arp{\\ts}226 = ``Atoms for Peace''); {\\it (c)}\\ IRAS{\\ts}19254--7245 (``The Super Antennae''); {\\it (d)}\\ IC{\\ts}4553/54 (Arp{\\ts}220).  The two at the top are LIGs whereas the two at the bottom are ULIGs.  Contours of H{\\ts}I 21-cm line column density ({\\it black}) are superimposed on deep optical ({\\it r}-band) images.  Inserts show a more detailed view in the {\\it K}-band (2.2{\\ts}$\\mu$m) of the nuclear regions of NGC{\\ts}4038/39, NGC{\\ts}7252, and IRAS{\\ts}19254--7245, and in the {\\it r}-band (0.65{\\ts}$\\mu$m) of Arp{\\ts}220.  White contours represent the CO(1$\\to$0) line integrated intensity as measured by the OVRO millimeter-wave interferometer.  No H{\\ts}I or CO interferometer data are available for the southern hemisphere object IRAS{\\ts}19254--7245. The scale bar represents 10{\\ts}kpc.} \\label{fig3} \\end{figure}  Although LIGs comprise the dominant population of extragalactic objects at ${\\it L}_{bol} > 10^{11}${\\ts}{\\it L}$_\\odot$, they are still relatively rare.  For example, Figure \\ref{fig1} suggests that only one object with ${\\it L}_{ir} >{\\ts}10^{12}${\\ts}{\\it L}$_\\odot$\\ will be found out to a redshift of $\\sim${\\ts}0.033, and indeed, Arp{\\ts}220 (${\\it z} = 0.018$) is the only ULIG within this volume. The total infrared luminosity from LIGs in the {\\it IRAS} Bright Galaxy Survey (BGS) is only $\\sim${\\ts}6\\% of the infrared emission in the local Universe \\cite{Soifer2}.   There are preliminary indications that ULIGs have been more numerous in the past.  Comparison of the space density of nearby ULIGs with the more distant population provides evidence for possible strong evolution in the luminosity function at the highest infrared luminosities.  Assuming pure density evolution of the form $\\Phi ({\\it z}) \\propto (1 + {\\it z})^{ n}$, \\cite{Kim} found ${\\it n} \\sim{\\ts}7\\pm${\\ts}3 for a complete flux-limited sample of ULIGs.   \\begin{figure}[p] \\begin{center} \\includegraphics[width=.8\\textwidth]{antennae_figgrey2.ps} \\end{center} \\caption[]{The upper figure from \\cite{Mirabel98} shows a superposition of the mid-infrared (12 -17 $\\mu$, countours) image of the Antennae galaxies obtained with the Infrared Space Observatory, on the composite optical image with V (5252 \\AA) and I (8269 \\AA) filters recovered from the Hubble Space Telescope archive .  About half of the mid-infrared emission from the gas and dust that is being heated by recently formed massive stars comes from an off-nuclear region that is clearly displaced from the most prominent dark lanes seen in the optical. The brightest mid-infrared emission comes from a region that is relativelly inconspicuous at optical wavelengths. The ISOCAM image was made with a 1.5$''$ pixel field of view. Contours are 0.4, 1, 3, 5, 10, and 15 mJy.  The lower figure shows the spectrum of the brightest mid-infrared knot and of the nuclei of NGC 4038 and NGC 4039. The rise of the continuum above 10 $\\mu$m and strong NeIII line emission observed in the brightest mid-infrared knot indicate that the most massive stars in this system of interacting galaxies are being formed in that optically obscured region, still enshrouded in large quantities of gas and dust.} \\label{fig4} \\end{figure}   The infrared properties for the complete {\\it IRAS} Bright Galaxy Sample have been summarized and combined with optical data to determine the relative luminosity output from galaxies in the local Universe at wavelengths $\\sim${\\ts}0.1--1000{\\ts}$\\mu m$ \\cite{Soifer2}.  Figure \\ref{fig2} illustrates how the shape of the mean spectral energy distribution (SED) varies for galaxies with increasing total infrared luminosity.  Systematic variations are observed in the mean infrared colors; the ratio ${\\it f}_{60}/{\\it f}_{100}$ increases while ${\\it f}_{12}/{\\it f}_{25}$ decreases with increasing infrared luminosity. Figure \\ref{fig2} also illustrates that the observed range of over 3 orders of magnitude in ${\\it L}_{ir}$\\ for infrared-selected galaxies is accompanied by less than a factor of 3--4 change in the optical luminosity.   \\cite{Sandersb} showed that a small but significant fraction of ULIGs, those with ``warm'' ({\\it f}$_{25}$/{\\it f}$_{60}$ $>${\\ts}0.3) infrared colors, have SEDs with mid-infrared emission ($\\sim$5--40{\\ts}$\\mu$m) over an order of magnitude stronger than the larger fraction of ``cooler'' ULIGs.  These warm galaxies (Figure \\ref{fig2} insert), which appear to span a wide range of classes of extragalactic objects including powerful radio galaxies (PRGs: ${\\it L}_{408MHz} \\geq 10^{25} W Hz^{-1}$) and optically selected QSOs, have been used as evidence for an evolutionary connection between ULIGs and QSOs (e.g. \\cite{Sandersa,Sandersb}).   There is a strong correlation between the broad band colors (from optical to far-infrared) and morphological type \\cite{SandersMirabel}. In particular, the fraction of objects that are interacting/merger systems appears to increase systematically with increasing infrared luminosity.  The imaging surveys of objects in the local universe \\cite{Sandersa,Melnick} have shown that the fraction of strongly interacting/merger systems increases from $\\sim${\\ts}10\\% at $log\\ ({\\it L}_{ir} / {\\it L}_{\\odot}$) = 10.5--11 to $\\sim${\\ts}100\\% at $log\\ ({\\it L}_{ir} / {\\it L}_{\\odot}$) $>${\\ts}12. In pannel (c) of Figure \\ref{fig3} is shown the \"Super-antennae\", which is the prototype of ULIG \\cite{Mirabel91}. ISO observations \\cite{Laurent99} have shown that more than 98\\% of the mid-infrared flux from this object comes from the southern component which hosts a Seyfert 2 nucleus.   From the detailed studies of nearby ultraluminous infrared galaxies the following conclusions were reached.  1) They are mergers of evolved gas-rich giant spiral galaxies (e.g. Milky Way with Andromeda), and not ``primival\" galaxies.  2) To boost the luminosity above 10$^{12}$ L$_{\\odot}$ the nuclei must have approached at least 10 kpc, namely, they are advanced mergers.  3) Due to the gravitational impact the interstellar gas decouples from the stars and large amounts of interstellar matter fall at high rates to the central region.  This is the condition to produce a nuclear starburst, and/or feed a supermassive black hole at super-Eddington accretion rates. To produce such large accretion rates, the gravitational potential wheels of massive buldges are needed.   A workshop on the question concerning the ultimate source of energy (starbursts versus AGN's) took place in Ringberg on October 1998. Below 2 10$^{12}$ L$_{\\odot}$ starbursts dominate the energy budget, but above 3 10$^{12}$ L$_{\\odot}$ AGN's seem to be always present and become an important source of energy. In this respect it is interesting to note that it is found with ISO that in the prototype Seyfert 2 galaxy NGC 1068, about 80\\% of the mid-infrared flux between 4 and 18 $\\mu$m comes from the AGN \\cite{LeFloc}.   A caveat for the subject of this conference is that the pre-encounter objects that merged at high redshifts must have been different from the metal-rich evolved galaxies merging at present. Another caveat is that ultraluminous IR galaxies at high redshifts may be very difficult to detect using the Lyman break technique. Due to the large amounts of dust in ultraluminous objects, very little or none continuum leaks out at ultraviolet wavelengths. Therefore, surveys with submillimeter arrays as ALMA will be needed to detect ultraluminous galaxies at high redshifts. ",
        "conclusions": "1) Scenarios on the history of star formation that use only observations in the UV and optical rest-frames result in luminosity functions that are strongly biased in the high luminosity end.   2) The most luminous nuclear and off-nuclear starbursts are enshrouded in dust.  In merging galaxies ISO revealed off-nuclear starburst knots with sizes $\\leq$100 pc that produce bolometric luminosities of up to 2 10$^{11}$ L$_{\\odot}$ (e.g. NGC 3690).  A few of these starburst knots can produce the overall bolometric luminosity of an ultraluminous galaxy such as Arp 220.   3) The observation with ISO of the nearest AGN to Earth (Centaurus A) opens the general question on whether the hosts of giant radio galaxies are symbiotic galaxies composed of spirals at the centre of giant ellipticals.   4) Mergers of disks can produce metal-rich elliptical galaxy cores.   5) Collisions between giant disk galaxies trigger the formation of dwarf galaxies out tidal debris. A fraction of these re-cycled galaxies become detached systems with diverse morphologies: blue compact dwarfs, dwarf irregulars, and irregulars of Magellanic type.   \\vskip .1in {\\it Acknowledgements:} Most of the work review here was carried out in collaboration with  D.B. Sanders, P-A. Duc,  V. Charmandaris and O. Laurent."
    },
    "0201/astro-ph0201530_arXiv.txt": {
        "abstract": "The XMM-Newton Reflection Grating Spectrometer Team has obtained observations of a large number of coronal sources of various activity levels, ages, and spectral types. In particular, X-ray bright RS CVn binary systems display saturated coronal emission with spectral lines characteristic of hot (10-30 MK) plasma. Furthermore, we have obtained XMM-Newton data from young solar analogs both within and outside the X-ray saturation reg\\-ime. We have simultaneously analyzed the EPIC MOS and RGS data from these objects and have obtained coronal abundances of various elements (e.g., C, N, O, Ne, Mg, Si, Fe). We show that there is evidence for a transition from an Inverse First Ionization Potential (FIP) effect in most active stars to a ``normal'' solar-like FIP effect in less  active stars. We discuss this result with regard to photospheric abundances. ",
        "introduction": "High-resolution X-ray spectra of stellar coronae obtained by \\textit{XMM-Newton} and \\textit{Chandra} now allow us to study in detail the rich forest of X-ray lines emitted by elements abundant in stellar coronae, such as C, N, O, Ne, Mg, Si, S, Ar, Ca, Fe, and Ni. In the past, stellar coronal abundances have frequently been determined using the moderate spectral resolution of CCD spectra from \\textit{ASCA} (e.g., \\cite{maudard-B1:drake96,maudard-B1:guedel99}) or from the low sensitivity spectrometers onboard \\textit{EUVE} (e.g., \\cite{maudard-B1:drake95,maudard-B1:laming96,maudard-B1:schmitt96,maudard-B1:drake97}). The abundance pattern in stellar coronae is complementary to the abundance pattern in the Sun: the solar corona, the solar wind, and solar energetic particles (and probably also galactic cosmic rays) display a so-called ``First Ionization Potential'' (FIP) effect, for which abundances of low-FIP ($<10$~eV) elements are enhanced relative to their respective photospheric abundances, while the abundances of high-FIP ($>10$~eV) elements are photospheric (\\cite{maudard-B1:feldman92}, see also \\cite{maudard-B1:meyer85}). Stellar coronal observations however often showed a deficiency of metals relative to the solar photospheric abundances (\\cite{maudard-B1:schmitt96}). \\textit{EUVE} spectra either indicated the absence of any FIP-related bias (\\cite{maudard-B1:drake95}), or a solar-like FIP effect (\\cite{maudard-B1:drake97}) in inactive stellar coronae. The new X-ray observatories \\textit{XMM-Newton} and \\textit{Chandra} combine the high spectral resolution with moderate effective areas to routinely obtain data allowing to measure the abundances in stellar coronae.  Recently, \\cite*{maudard-B1:brinkman01} showed a trend towards enhanced high-FIP elemental abundances, while low-FIP abundances are depleted; this effect was dubbed the ``Inverse FIP'' (IFIP) effect. Other active stars showed a similar trend (\\cite{maudard-B1:guedel01a,maudard-B1:guedel01b}), except the intermediately active Capella (\\cite{maudard-B1:audard01a}). Note however that stellar coronal abundances have often been normalized to the \\emph{solar} photospheric abundances, while they should better be normalized to the \\emph{stellar} photospheric abundances. The latter are difficult to measure. Nevertheless, for some stars, photospheric abundances are known. The uncertainty introduced by photospheric  abundances can then be removed. We will show that there is a transition from an IFIP to a normal FIP effect in the long-term evolution of the coronae from active to inactive solar analogs. We will use data of bright active RS CVn binary systems to complement the study as well. Finally, the variation of stellar coronal abundances during flares will be discussed. ",
        "conclusions": "High-resolution X-ray spectra of magnetically active stars have been investigated with the Reflection Grating Spectrometers on board \\textit{XMM-Newton}. The high-energy data ($> 1.5$~keV) of the EPIC CCD spectra were used to better constrain the high-temperature part of the emission measure distributions and to profit from the presence of H-like and He-like transitions of Si, S, Ar, and Ca. It was found that the most active stars, such as the bright RS CVn binary systems, show a marked depletion of low-FIP elements (e.g., Fe, Mg, Si) relative to high-FIP elements (e.g., C, N, O, Ne), opposite to the FIP effect observed in the solar corona. This ``inverse FIP'' effect is however not observed in the intermediately active RS CVn binary Capella. Since their photospheric abundances are mostly unknown or not reliable, one can hypothesize that the observed FIP bias is simply a reflection of their photospheric composition.  To remove the uncertainty of surface abundances, we have analyzed high-resolution X-ray spectra of solar ana\\-logs of known photospheric composition (close to solar). These solar-like stars span a wider range of coronal activity (from inactive to active) and represent the evolution of the solar corona in time. We have found a transition from a depletion of low-FIP elements (relative to high-FIP elements) in the most active stars toward a marked enhancement of their abundances in the inactive stars. On the other hand, the abundances of high-FIP elements do not vary with the activity level (or coronal temperature), relative to O. The IFIP effect found in the active RS CVn binary systems fit well into this transition, under the assumption that their photospheric composition is also close to solar. Similarly, the solar FIP effect (enhancement of low-FIP elements by factors of 4--8) fits into this picture. However, although the scenario of correlating the activity level with the FIP bias is tempting, it may be too simplistic; indeed, such scenario does not explain the absence of any FIP bias in the corona of the old, inactive Procyon (\\cite{maudard-B1:drake95}).  We have put forward first ideas to explain the inverse FIP effect seen in active stars: downward propagating electrons detected by their gyrosynchrotron emission in active stars could prevent chromospheric ions (mostly low-FIP elements) from escaping into the corona by building up a downward-pointing electric field (\\cite{maudard-B1:guedel02}). As the density of high-energy electrons decreases with decreasing activity, the inverse FIP effect is quenched. During large flares, however, the high-energy electrons heat a significant portion of the chromosphere to bring up a near-photospheric mixture of elements into the corona, and this  effect has indeed been observed (\\cite{maudard-B1:guedel99,maudard-B1:audard01b}). The new results by \\textit{XMM-Newton} and \\textit{Chandra} have opened a new field of research relevant to the physics of heating and dynamics of outer stellar atmospheres.    \\begin{figure*}[!ht] \\begin{center} \\includegraphics{maudard-B1_fig1.eps}\t% \\end{center} \\caption{RGS spectra of bright active RS CVn binary systems. The sources have been ordered with decreasing activity levels (or average coronal temperature) from top to bottom. Major emission lines have been labeled.} \\label{maudard-B1_fig:fig1} \\end{figure*}  \\begin{figure*}[!ht] \\begin{center} \\includegraphics{maudard-B1_fig2.eps}\t% \\end{center} \\caption{Coronal abundance normalized to oxygen in RS CVn binaries as a function of the First Ionization Potential. Solar photospheric abundances from Anders \\& Grevesse (1989) were used, except for {\\rm Fe} (Grevesse \\& Sauval 1999).} \\label{maudard-B1_fig:fig2} \\end{figure*}   \\begin{figure*}[!ht] \\begin{center} \\includegraphics{maudard-B1_fig3.eps}\t% \\end{center} \\caption{RGS spectra of solar analogs. Their order is set similarly to Fig.~\\ref{maudard-B1_fig:fig1}. The arrows designate lines with similar maximum formation temperature; hence different line ratios of the {\\rm Fe}~\\textsc{xvii} line at 15\\AA\\  and the {\\rm Ne}~\\textsc{ix} line suggest differences in coronal abundances in each star.} \\label{maudard-B1_fig:fig3} \\end{figure*}  \\begin{figure*}[!ht] \\begin{center} \\includegraphics{maudard-B1_fig4.eps}\t% \\end{center} \\caption{Coronal abundance normalized to oxygen in solar analogs as a function of the First Ionization Potential. Note that the activity level decreases (their age increases) from top to bottom. Similar photospheric abundances have been taken as in Fig.~\\ref{maudard-B1_fig:fig2}.} \\label{maudard-B1_fig:fig4} \\end{figure*}   \\begin{figure*}[!ht] \\begin{center} \\includegraphics[width=\\textwidth]{maudard-B1_fig5.eps}\t% \\end{center} \\caption{Coronal abundances (normalized to {\\rm O}) as a function of the average coronal temperatures, for {\\rm Fe} (low-FIP; left) and {\\rm Ne} (high-FIP; right). The data include solar analogs and RS CVn binaries. } \\label{maudard-B1_fig:fig5} \\end{figure*}  \\begin{figure*}[!ht] \\begin{center} \\includegraphics[width=\\textwidth]{maudard-B1_fig6.eps}\t% \\end{center} \\caption{Coronal abundances (normalized to {\\rm O}) as a function of the average coronal temperature during a large flare in HR~1099. Left panel gives {\\rm Fe/O} ratios, while the right panel gives {\\rm Ne/O} ratios. `Q' stands for quiescent, `R' for flare rise, and `P' for flare peak.} \\label{maudard-B1_fig:fig6} \\end{figure*}"
    },
    "0201/astro-ph0201277_arXiv.txt": {
        "abstract": "We give an explanation for the origin of various properties observed in local infrared galaxies, and make predictions for galaxy counts and cosmic background radiation (CBR), by a new model extended from that for optical/near-infrared galaxies. Important new characteristics of this study are that (1) mass scale dependence of dust extinction is introduced based on the size-luminosity relation of optical galaxies, and that (2) the big grain dust temperature $T_{\\rm dust}$ is calculated based on a physical consideration for energy balance, rather than using the empirical relation between $T_{\\rm dust}$ and total infrared luminosity $L_{\\rm IR}$ found in local galaxies, which has been employed in most of previous works. Consequently, the local properties of infrared galaxies, i.e., optical/infrared luminosity ratios, $L_{\\rm IR}$-$T_{\\rm dust}$ correlation, and infrared luminosity function are outputs predicted by the model, while these have been inputs in a number of previous models. Our model indeed reproduces these local properties reasonably well. Then we make predictions for faint infrared counts (in 15, 60, 90, 170, 450, and 850 $\\mu$m) and CBR by this model. We found considerably different results from most of previous works based on the empirical $L_{\\rm IR}$-$T_{\\rm dust}$ relation; especially, it is shown that the dust temperature of starbursting primordial elliptical galaxies is expected to be very high (40--80K), as often seen in starburst galaxies or ultra luminous infrared galaxies in the local and high-$z$ universe. This indicates that intense starbursts of forming elliptical galaxies should have occurred at $z \\sim 2$--3, in contrast to the previous results that significant starbursts beyond $z \\sim 1$ tend to overproduce the far-infrared (FIR) CBR detected by {\\sl COBE}/FIRAS. On the other hand, our model predicts that the mid-infrared (MIR) flux from warm/nonequilibrium dust is relatively weak in such galaxies making FIR CBR, and this effect reconciles the {\\it prima facie} conflict between the upper limit on MIR CBR from TeV gamma-ray observations and the {\\sl COBE}\\ detections of FIR CBR. The intergalactic optical depth of TeV gamma-rays based on our model is also presented. ",
        "introduction": "{}To understand when and how stars and galaxies formed in the universe is one of the most fundamental issues in modern astronomy. The energy emitted from stars as a result of nuclear fusion is radiated in two modes: direct stellar emission ranging from optical to near-infrared (NIR) wavelengths, and emission in mid- and far-infrared (MIR and FIR) wavelengths from dust particles heated by stellar radiation field.\\footnote{In this paper, the term `infrared' refers to the emission from dust particles in the MIR and FIR including submillimeter, but excluding direct stellar emission in the NIR.} Faint galaxy counts and the cosmic background radiation (CBR), as well as the properties of galaxies in the local universe, give us important clues to understand galaxy formation.  In optical/NIR wavelengths, the sensitivity of existing telescopes now reaches a depth sufficient to resolve more than 80--90\\% of CBR from galaxies (Totani et al. 2001a), thanks to very deep optical surveys such as the Hubble Deep Fields (HDFs: Williams et al. 1996; Williams et al. 2000) and the NIR surveys such as the Subaru Deep Field (SDF: Maihara et al. 2001). There are a number of theoretical models in various approaches to be compared with these data, and in fact these data can be reasonably explained in the framework of the Big-Bang cosmology and structure formation induced by the cold dark matter (CDM). In contrast, infrared observations have not yet reached such depth and star formation activity hidden by dust extinction is still poorly known, despite the dramatic progress of observations achieved by satellites such as {\\sl IRAS} (e.g., Soifer et al.\\ 1987 and references therein) and {\\sl ISO} (e.g.\\ Puget et~al.\\ 1999; Oliver et~al.\\ 2000; Okuda~2000; for a review, see e.g., Genzel \\& Cesarsky 2000). The SCUBA of JCMT (Holland et al.\\ 1999) has opened a new window of submillimeter wavelengths to probe the dusty galaxies at high-$z$, and it seems that a considerable part of CBR from submm galaxies has been resolved (e.g., Smail, Ivison, \\& Blain 1997; Hughes et al.\\ 1998; Blain et~al.\\ 1999; Blain et al.\\ 2000; Eales et~al.\\ 1999, 2000; and Barger et al.\\ 1998, 1999). However, insufficient angular resolution does not allow us to identify most of these objects in other wavelengths. Forthcoming projects such as {\\sl SIRTF}, {\\sl ASTRO-F}, SOFIA, {\\sl Herschel Space Observatory}, {\\sl NGST}, and ALMA would bring about revolutionary progress from this situation. Comparison of these data with theoretical models of infrared galaxies will be the key to understand the ``dark side'' of galaxy formation in the next decade of extragalactic astronomy.  Theoretical modeling of galaxy formation and evolution has a long history. In optical/NIR wavelengths, most theoretical modelings can be roughly divided into two categories: the so-called ``backwards'' approach and ``{\\it ab initio}'' approach. In the former approach (Tinsley 1980; Yoshii \\& Takahara 1988; Fukugita et al. 1990; Rocca-Volmerange \\& Guiderdoni 1990; Yoshii \\& Peterson 1991, 1995; Pozzetti et al. 1996, 1998; Jimenez \\& Kashlinsky 1999; Totani \\& Yoshii 2000), the luminosity function of local galaxies is used as an input to normalize the number density of galaxies. The local properties of galaxies such as multi-band colors and chemical properties are also used to construct a reasonable model of star formation history and luminosity evolution of galaxies based on the stellar population synthesis method. The evolution is then probed backwards into the past to predict observables such as galaxy counts and redshift distributions. The formation epoch and merging history of galaxies cannot be predicted in this framework, and hence they are introduced as phenomenological parameters that can be inferred from comparison with observational data. In the latter approach (Kauffmann et al. 1993; Cole et al. 1994, 2000; Somerville \\& Primack 1999; Nagashima et al. 2001), on the other hand, the formation epoch and merging history of galaxies are predicted by the standard theory of structure formation in the CDM universe. In these models the local properties such as luminosity function are outputs of the model which should be compared with observations. However, although the formation and evolution of dark matter halos are rather well understood and can be properly predicted, our knowledge about baryonic processes such as star formation, supernova feedback, or galaxy merging is still very poor, and a number of phenomenological parameters must be incorporated, making the comparison of a model and observed data rather complicated.  Both of these two approaches have been used also to make predictions of infrared galaxy counts and infrared CBR (see Franceschini et al. 1994 for the former, Guiderdoni et al. 1998 and Devriendt \\& Guiderdoni 2000 for the latter, and Tan, Silk, \\& Balland 1999 for a somewhat hybrid approach between the two). However, because of our poor knowledge and understanding of dust formation and emission, there are considerable difficulties in theoretically predicting infrared luminosity, spectral energy distribution (SED), and their evolution of dust emission, compared with the optical/NIR modeling based on the stellar population synthesis method. This is the motivation of another ``empirical'' approach\\footnote{In several publications this approach is also referred to a ``backwards'' approach, but here we discriminate between the ``backwards'' approach mentioned earlier and the ``empirical'' approach.} in the MIR and FIR wavelengths. In this kind of approach, the infrared luminosity evolution is introduced by some phenomenological functional forms and constraints on luminosity evolution is derived from comparison with observed data (Beichman \\& Helou 1991; Oliver, Rowan-Robinson \\& Saunders 1992; Blain \\& Longair 1993, 1996; Pearson \\& Rowan-Robinson 1996; Malkan \\& Stecker 1998, 2001; Takeuchi et al. 1999, 2001a, 2001b; Roche \\& Eales 1999; Gispert, Lagache, \\& Puget 2000; Wang \\& Biermann 2000; Xu et al. 2001; Franceschini et al. 2001).  There is yet another approach so-called ``cosmic chemical evolution'', in which the universe is treated as a uniform medium and the mean gas consumption by stars in the whole universe is considered (Pei \\& Fall 1995; Fall, Charlot, \\& Pei 1996; Pei, Fall, \\& Hauser 1999; Sadat, Guiderdoni, \\& Silk 2001). The cosmic star formation history is determined by the redshift evolution of the cosmic mass density of neutral gas and the cosmic mean metallicity, which is inferred from quasar absorption line systems. Although this is a simple and beautiful approach to predict CBR (but not galaxy counts),  several problems exist in this kind of approach, which could be serious especially in modeling infrared radiation from dust. First, quasar absorption systems are inevitably biased to gas-rich systems, i.e., systems in which star formation has been relatively inefficient, and hence their mean metallicity does not necessarily represent the mean of the whole universe. For example, if an elliptical galaxy has formed at $z \\sim 4$, this galaxy has been already in passive evolution phase at $z \\sim 2$. Then it can never be observed as a quasar absorption system because of complete exhaustion of interstellar gas, although significant stars and metals have been produced. Galaxies in dusty starburst phase may not be traced by absorption line systems, either, because background quasars may be completely extincted and cannot be observed. Therefore, the metallicity and star formation activity observed in absorption line systems must be an underestimate of the true mean in the universe. Furthermore, extinction by interstellar dust and reradiation in infrared bands are very sensitively dependent on dust opacity and its geometrical distribution within a galaxy, but this important information is completely missed.  In this paper we construct a model of infrared galaxy counts and CBR based on the backwards approach for luminosity evolution of galaxies, by extending the model for optical/NIR galaxy counts (Totani \\& Yoshii 2000, hereafter TY00) to include dust emission component in infrared bands. After the work by Franceschini et al. (1994, F94 hereafter), only a limited number of models have been published based on this approach covering FIR bands, while there has been a dramatic increase of quality and quantity of observational data in the local universe as well as for faint/high-$z$ galaxies. Although the backwards approach has a disadvantage that formation epoch and merger history must be treated phenomenologically, it has an advantage that the number of model parameters is much fewer, and connection of the model galaxies to the local galaxy populations is clearer, than the {\\it ab initio} approach. Another reason for the relatively low activity of work in this direction is the difficulty of constructing realistic evolutionary models that can consistently describe both the optical and infrared data at the local universe as well as high-$z$, as argued by Franceschini et al.\\ (2001). This is indeed what motivated more empirical studies confined to the infrared bands.  The aim of this paper is to explain the bulk of optical and infrared data consistently by realistic evolutionary models of known and relatively normal galaxy populations at the local universe, without introducing hypothetical populations or parametric description of luminosity evolution. An important improvement of our model from the F94 model is that mass scale dependence is introduced for dust extinction and reradiation, based on the size-luminosity relation of galaxies observed in the optical bands. In the F94 model, there was no physical difference in model galaxies of a given type with different masses; massive galaxies were simply a scaled-up version of smaller objects. We will show that this point is essential to understand various local properties of infrared galaxies.  Another important characteristic of this work is modeling of the infrared SED evolution. In most of the previous infrared models, including {\\it ab initio} and empirical approaches, the dust SED is empirically modeled in such a way that the dust SED is a one-parameter family specified by the total infrared luminosity of galaxies. The relation between the dust SED and luminosity, which is observed in local infrared galaxies, is characterized by a gradual increase of characteristic dust temperature $T_{\\rm dust}$ with infrared luminosity $L_{\\rm IR}$ and is often expressed simply as a linear relation between $T_{\\rm dust}$ and $\\log L_{\\rm IR}$ (e.g.\\ Smith et al. 1987; Soifer \\& Neugebauer 1991). However in this paper, we argue that this relation is not physically warranted for high-$z$ galaxies. In a physical sense, the luminosity is {\\it extensive}, i.e., a quantity which is dependent upon the amount of substance present in the system. On the other hand, the dust temperature is {\\it intensive}, which is a specific characteristic of the system and is independent of the amount of material concerned. Therefore these two must not be related by only one sigle relation. At least there should be another extensive quantity such as the total mass of dust in a galaxy $M_{\\rm dust}$. In the empirical $L_{\\rm IR}$-$T_{\\rm dust}$ relation, the extensive scale of an object is completely missed. The empirical $L_{\\rm IR}$-$T_{\\rm dust}$ relation of local infrared galaxies should be considered as a projection of the plane between $L_{\\rm IR}$, $T_{\\rm dust}$, and $M_{\\rm dust}$. Here we will construct a model on physical basis to describe the infrared luminosity as a two-parameter family of $M_{\\rm dust}$ and $T_{\\rm dust}$. The dust temperature is physically calculated, without using the empirical $L_{\\rm IR}$-$T_{\\rm dust}$ relation. In addition to the thermal emission from big grain dust, the warm/nonequilibrium components of dust emission from small grain dusts and PAH features are also added to the model SED to be consistent with the observed infrared SED of galaxies.  It is important to check that the model can successfully reproduce the local properties of infrared galaxies, before predicting the high-$z$ quantities such as counts and CBR, since in our model all the inputs are the optical/NIR properties of local galaxies, and those in infrared bands are outputs predicted by the model. We will compare the model prediction with the observed properties of local infrared galaxies, such as the correlation between optical and infrared luminosity, the $L_{\\rm IR}$-$T_{\\rm dust}$ relation, and infrared luminosity function. Indeed, we will show that our model gives a reasonable explanation for these relations as well as the scatter along the mean relation. Recently Granato et al.\\ (2000) have presented a very sophisticated semi-analytic model of infrared galaxy formation to calculate various properties of local infrared galaxies, although predictions for high-$z$ galaxies have not been made yet. They have also found that most of the observed properties can be reproduced in their model framework. However, a large number of parameters have been introduced to treat various physical processes, and the origin of the observed properties of local infrared galaxies has not been clearly discussed. Here we try to shed light on the underlying physics of these properties by a model with a much fewer number of adjustable parameters.  Then we will make predictions of infrared galaxy counts and CBR by our backwards evolution model. We found that the empirical $L_{\\rm IR}$-$T_{\\rm dust}$ relation breaks down at high-$z$, especially in the intense starbursts expected at the formation of elliptical galaxies. The dust temperature should be much higher (40--80K) than predicted by mere extrapolation of the empirical $L_{\\rm IR}$-$T_{\\rm dust}$ relation. We will discuss the implications of our results for the star formation history of the universe, especially the formation epoch of elliptical galaxies. Our result also has an interesting implication for the {\\it prima facie} conflict (Stecker 2000) between the upper limit on MIR CBR from TeV gamma-ray observations and the FIR CBR detections by the {\\sl COBE}\\ satellite. The calculation of intergalactic optical depth of TeV gamma-rays based on our model will also be presented.  The paper will be organized as follows. The model of galaxy evolution in the optical bands, on which the infrared modeling is based, will be summarized in \\S \\ref{section:model-optical}. Then the extension of the optical/NIR model to include the dust emission is described in detail in \\S \\ref{section:model-infrared}. Our model predicts local properties of infrared galaxies, and they are compared with the observations in \\S \\ref{section:local}. The predictions of faint infrared counts and CBR will be presented to be compared with observed data, in \\S \\ref{section:high-z}. Implications and discussions on our results are given for several topics in \\S \\ref{section:discussion}, and finally we summarize and conclude this paper in \\S \\ref{section:conclusions}. Unless otherwise stated, we will adopt a cosmological model with $h = H_0 / (100\\rm km/s/Mpc) = 0.7$, $\\Omega_0 = 0.2$, $\\Omega_\\Lambda=0.8$. ",
        "conclusions": "\\label{section:conclusions} In this paper we developed a new model of counts and cosmic background radiation (CBR) of infrared galaxies observed by emission from heated interstellar dust, by extending a model for optical/NIR galaxies (Totani \\& Yoshii 2000). Five morphological types of galaxies (E/S0, Sab, Sbc, Scd, and Sdm) are taken into account and their number densities are normalized by type-dependent $B$-band luminosity function at the local universe. Their luminosity evolution is traced backwards in time based on star formation histories inferred from the present-day optical/NIR SEDs and chemical properties. Formation epoch of galaxies is a parameter for which we tried a redshift range of $2 \\leq z_F \\leq 5$. Pure luminosity evolution without number evolution is assumed in our baseline model, but some number evolution is also tested in a phenomenological way. The model has already been compared comprehensively with the counts and redshift distributions of galaxies observed in optical and NIR bands, and found to be in reasonable agreement with the data (Totani \\& Yoshii 2000; Totani et al. 2001c).  Relatively rare populations of AGNs and ULIRGs seen in the local universe, whose contribution to counts and CBR is expected to be small, are not included in our model, and we tried to explain the bulk of infrared data by normal galaxy populations at the local universe and their ancestors at high redshifts. On the other hand, our analysis strongly indicates that the primordial elliptical galaxies are very similar to dusty starburst galaxies or ULIRGs with very high dust temperature.  There are two important new characteristics in our model that are different from previous models of infrared galaxies: (1) mass scale dependence of dust extinction is introduced by the observed size-luminosity relation of optical galaxies, and (2) dust temperature is determined by physical consideration of energy balance, rather than using the empirical relation between the total infrared luminosity ($L_{\\rm IR}$) and characteristic dust temperature ($T_{\\rm dust}$) of local galaxies that has been used in a number of previous models. As a result, the local properties of infrared galaxies, such as optical/infrared luminosity ratios, correlation between infrared luminosity and dust temperature, and infrared luminosity function, are outputs that should be compared with observed data. Indeed we found that our model quantitatively reproduces the observed infrared properties at the local universe. The key to understand these scaling properties is the size-luminosity relation of galaxies; surface brightness and dust column density increase with increasing optical galaxy luminosity, and hence more massive galaxies should be more dusty. This gives a quantitative explanation for the observed correlation between optical and infrared luminosities. Furthermore, massive galaxies should emit more energy as dust emission per unit mass of dust, and hence the energy balance inevitably results in higher temperature for larger galaxies, as observed for local infrared galaxies. The scatter along the mean $L_{B}$-$r_e$ relation is comparable with those in $L_{\\rm IR}$-$L_{B}$ and $L_{\\rm IR}$-$T_{\\rm dust}$ relations. These effects result in much faster increase of 60$\\mu$m luminosity when $L_B$ is increased, and giving an explanation for the much broader shape of 60$\\mu$m luminosity function than the Schechter function of the optical luminosity function.  Then we predicted faint source counts and CBR composed of high-$z$ galaxies. Our baseline model assumes a cosmological model with ($h, \\Omega_0, \\Omega_\\Lambda$) = (0.7, 0.2, 0.8), pure luminosity evolution after formation at $z_F = 3$, and screen distribution of dust. We found that this baseline model is in reasonable agreement with all available data of galaxy counts in six wavebands (15, 60, 90, 170, 450, and 850$\\mu$m) and CBR. Therefore our model, based only on present-day normal galaxy populations and their evolution, reasonably fits to all available data from optical to submillimeter wavebands, though some modest number evolution may be required for late-type galaxies. The high-$z$ starburst galaxies discovered by SCUBA are quantitatively well explained in our model by the emergence of starbursts at the formation of present-day elliptical galaxies at $z \\sim 3$.  We also tested that slab-type dust distribution as well as the screen dust used in the baseline model, and found that the screen model gives much better fit to the observed data of the local infrared luminosity function, galaxy counts, and CBR. Although a pure screen distribution may seem unlikely, it is rather reasonable to expect that a part of dust particles behave like an effective screen, because of stellar/galactic wind or inhomogeneity of interstellar medium. If it is the case, and the dust opacity to optical/NIR is much larger than the unity, the screen model should be better than the slab-type prescription. It is also supported by extremely red colors of local starburst galaxies (Calzetti, Kinney, \\& Storchi-Bergmann 1994) or recently discovered hyper extremely red objects (Totani et al. 2001b), which cannot be explained simply by slab-type dust. Therefore we conclude that screen dust is a better phenomenological description than slab dust at least in a study of this kind.  The most drastic difference of our model from previous ones is that the dust temperature of starbursting elliptical galaxies is predicted to be much higher ($\\sim$ 40--80K) than that extrapolated by the empirical $L_{\\rm IR}$-$T_{\\rm dust}$ relation of local infrared galaxies. This is because starbursting elliptical galaxies should emit much larger amount of energy as dust emission {\\it per unit dust mass}, than local galaxy populations. On the other hand, such high temperature is similar to those found in local ULIRGs or high-$z$ dust starbursts observed in sub-mm bands. Thus, our result gives a further support to an idea that the progenitor of present-day elliptical galaxies or bulges are dusty starbursts.  There is an important implication for the cosmic star formation history, which is very different from previous results. A number of papers based on the empirical $L_{\\rm IR}$-$T_{\\rm dust}$ relation claimed that cosmic star formation rate beyond $z \\gtrsim$ 1 must turn over and keep constant or decline, otherwise it would produce too much submillimeter CBR compared with the {\\sl COBE}/FIRAS data, by the redshifted dust emission from high-$z$ galaxies (e.g., Gispert, Lagache, \\& Puget 2000; Takeuchi et al.\\ 2001a; Franceschini et al.\\ 2001). However, in our model the dust temperature of forming elliptical galaxies is much higher than in the previous models. As a result, although our model assumes very strong starbursts at $z_F \\sim$ 3 in primordial elliptical galaxies and the cosmic SFR at $z_F \\gtrsim$ 2--3 is even higher than the peak at $z \\sim$ 1 suggested by optical observations, the prediction is in good agreement not only with the {\\sl COBE}/FIRAS CBR measurements but also with the counts and redshift estimation for submillimeter sources revealed by SCUBA. In our model the dusty starbursts in primordial elliptical galaxies hardly contribute to the cosmic SFR measured by optical observations.  Another result that is significantly different from previous studies is the smaller ratio of MIR/FIR CBR flux compared with models based on the empirical $L_{\\rm IR}$-$T_{\\rm dust}$ relation of local infrared galaxies (e.g., Stecker 2000). It should be noted that our model does include the warm/nonequilibrium components of dust emission that dominate the thermal emission from big grain dust at the MIR region, while several previous models gave very low MIR CBR flux as a valley in the CBR spectrum simply because they did not take them into account (e.g., MacMinn \\& Primack 1996; Fall, Charlot, \\& Pei 1996). The reason we get this result is the trend that the warm/nonequilibrium components become less significant compared with thermal big-grain dust emission, with increasing temperature of dust. This is inferred from the infrared spectrum of local galaxies, and it is also expected from physical consideration. If this trend is taken for granted for high-$z$ galaxies, the MIR/FIR flux ratio of dust emission from forming elliptical galaxies should be much smaller than that of local infrared galaxies, since the dust temperature of starbursting elliptical galaxies is found to be very high in our model. Therefore these galaxies dominate in the FIR peak of CBR while they have very small contribution in MIR CBR. This effect is significant enough to resolve the {\\it prima facie} conflict (Stecker 2000) between the upper limits on MIR CBR from TeV gamma-ray observations and FIR CBR detections by {\\sl COBE}/DIRBE and FIRAS.  The authors have financially been supported by the JSPS Fellowship. The authors thank Hirohisa Nagata for providing their temperature data of {\\sl IRAS} galaxies."
    },
    "0201/astro-ph0201088_arXiv.txt": {
        "abstract": "{ We present optical photometry, spectroscopy and photopolarimetry, as well as {\\em ASCA} X-ray observations, of the recently discovered intermediate polar 1WGA J1958.2+3232. Through the first detection of an optical beat frequency, we confirm the previously tentative suggestion that the spin period of the white dwarf is twice the X-ray and optical pulsation period, which we also confirm in each case. We detect an orbital modulation in each of the U, B, V, R and I bands for the first time, and suggest that the true orbital period is the --1d alias of that previously suggested. We also confirm the presence of circular polarization in this system, detecting a variable polarization which has opposite signs in each of the B and R bands. The double peaked pulse profile and oppositely signed polarization pulses suggest that 1WGA J1958.2+3232 accretes onto both magnetic poles via a disc which is truncated relatively close to the white dwarf. ",
        "introduction": "Intermediate polars (IPs) are semi-detached interacting binaries in which a magnetic white dwarf accretes material from a Roche-lobe filling, usually late-type, main sequence companion star. The accretion flow from the secondary proceeds towards the white dwarf either through an accretion disc, an accretion stream, or some combination of both (known as disc overflow accretion), until it reaches the magnetospheric radius. Here the material attaches to the magnetic field lines and follows them towards the magnetic poles of the white dwarf. The infalling material that originates from an accretion disc takes the form of arc-shaped accretion curtains, standing above the white dwarf surface. At some distance from this surface, the accretion flow undergoes a strong shock, below which material settles onto the white dwarf, releasing X-ray to optical emission. Since the magnetic axis is offset from the spin axis of the white dwarf, this gives rise to the defining characteristic of the class, namely X-ray (and usually optical) emission pulsed at the white dwarf spin period. Additionally, as the X-ray `beam' sweeps around the system, there is the possibility that some fraction of the emission will be reprocessed from structures such as the companion star or a bulge at the edge of an accretion disc. This will give rise to further optical emission pulsed at the lower orbital sideband of the spin frequency, namely the spin frequency of the white dwarf in the reference frame of the binary. Several IPs show a dominant optical pulsation at this orbital sideband frequency. X-ray pulsation at the orbital sideband frequency is due to an intrinsic modulation arising from pole-switching in the case of stream-fed accretion. Several IPs show this type of modulation (at least at some epochs) too. About twenty confirmed intermediate polars are now recognised with a similar number of candidate systems having been proposed. Comprehensive reviews of various aspects of their behaviour are given by Patterson (\\cite{Patt}), Warner (\\cite{War95}), Hellier (\\cite{Hell95}; \\cite{Hell96}) and Norton (\\cite{Nor95}). ",
        "conclusions": "Optical photometry from the JKT has allowed us to confirm the detection of an optical pulsation in 1WGA J1958.2+3232 at a period of 733.33~s and define a refined ephemeris. {\\em ASCA} observations also confirm the presence of an X-ray pulsation at the same period as seen in the optical. Our detection of a second optical pulsation period of 1587.24~s, identified with the beat period of the system, implies that the spin period of the white dwarf is {\\em twice} the short pulsation period, i.e.  $P_{\\rm spin} = 1466.66$~s and that the orbital period is the --1 day alias of the strongest low frequency peak in our power spectrum, i.e. $P_{\\rm orb} = 5.387$~h. This is also the --1 day alias of the previously reported orbital period. We detect an orbital modulation in each of the U, B, V, R and I bands for the first time.  We also confirm the presence of circularly polarized emission from this source, which is thus only the fifth IP to exhibit such behaviour. The circular polarization is negative in the R-band and positive in the B-band, and in each case shows evidence for variation across the spin cycle. The level of polarization in the R-band, at $\\sim -1\\%$ to $\\sim -5\\%$, is greater than  that seen by Uslenghi et al (2001) who saw a mean level of $\\sim -0.5\\%$, whilst our detection of positive circular polarization in the B-band is a first detection in that band.  Since the spin period of the white dwarf is twice the pulsation period observed in both optical and X-ray flux, the pulse profile is actually double peaked. Following Norton et al (\\cite{Nor99}), a double peaked X-ray pulse profile indicates that two-pole accretion in 1WGA J1958.2+3232 conspires to produce two peaks per revolution of the white dwarf. In this model, the white dwarf has a relatively low magnetic field strength, and consequently an accretion disc which is truncated only relatively close to the white dwarf. This leads to large footprints of the accretion curtains and an optical depth to X-rays across the accretion regions which is largest in a direction parallel to the white dwarf surface, and smallest perpendicular to the surface. (This is the reverse of the standard accretion curtain model which leads to a single peaked X-ray pulse profile.) Confirmation of this model would require pulse phase resolved optical spectroscopy in order to measure the direction of flow of the accreting material corresponding to the phases of pulse maximum and minimum. Since this model also predicts that the accretion flow is essentially via a disc, an X-ray beat modulation would not be expected to be present, in accord with all observations so far, including the {\\em ASCA} data presented here.  The different polarities of the circular polarization observed in the B- and R-band data, together with light curve variations, also suggest two pole accretion. We suggest that between spin phases 0.0 and 0.5 the negative polarization accretion pole is seen, and after phase 0.5 emission from this pole reduces. Between spin phase 0.3 and 0.8, the positive polarization accretion pole is seen in the B-band. The different colours of the two accretion poles may indicate that the poles have significantly different magnetic field strengths, possibly as a result of an offset dipole field structure. We note however, that the detection of circular polarization from an IP is generally taken to be an indication of a relatively strong magnetic field, in contradiction to the implication of the double-peaked pulse profile model.  There is though no evidence for a strong magnetic field in this system. The figure of $B = 8$~MG stated by Uslenghi et al (2001) is merely a value chosen for illustrative purposes and is in no way `derived' from their polarization measurements. We suggest that the magnetic field strength of 1WGA J1958.2+3232 is low enough to allow a small disc truncation radius, and consequently produce a double peaked X-ray pulse profile, but high enough that geometrical effects allow a detection of polarized emission. For instance, in a system with a more symmetrical field pattern than is implied here, the (positive) polarization from one pole may be effectively cancelled out by that (negative) from the other pole, so yielding a net polarization close to zero and hence undetectable. It may be that the only reason we detect polarization in 1WGA J1958.2+3232 is that the two poles have sufficiently different emission properties that a net polarization remains. In this interpretation, the question of whether each IP exhibits polarized emission is not only dependent upon the strength of the white dwarf magnetic field, but also on the geometry of the system and that of its magnetic field."
    },
    "0201/astro-ph0201041_arXiv.txt": {
        "abstract": "We calculate numerically the collapse of slowly rotating, non-magnetic, massive molecular clumps of masses 30~M\\sol, 60~M\\sol, and 120~M\\sol, which conceivably could lead to the formation of massive stars. Because radiative acceleration on dust grains plays a critical role in the clump's dynamical evolution, we have improved the module for continuum radiation transfer in an existing 2D (axial symmetry assumed) radiation hydrodynamic code.  In particular, rather than using ``grey'' dust opacities and ``grey'' radiation transfer, we calculate the dust's wavelength-dependent absorption and emission simultaneously with the radiation density at each wavelength and the equilibrium temperatures of three grain components: amorphous carbon particles, silicates, and ``dirty ice''-coated silicates. Because our simulations cannot spatially resolve the innermost regions of the molecular clump, however, we cannot distinguish between the formation of a dense central cluster or a single massive object.  Furthermore, we cannot exclude significant mass loss from the central object(s) which may interact with the inflow into the central grid cell.  Thus, with our basic assumption that all material in the innermost grid cell accretes onto a single object, we are only able to provide an upper limit to the mass of stars which could possibly be formed.  We introduce a semi-analytical scheme for augmenting existing evolutionary tracks of pre-main sequence protostars by including the effects of accretion.  By considering an open outermost boundary, an arbitrary amount of material could, in principal, be accreted onto this central star.  However, for the three cases considered (30~M\\sol, 60~M\\sol, and 120~M\\sol\\ originally within the computation grid), radiation acceleration limited the final masses to 31.6~M\\sol, 33.6~M\\sol, and 42.9~M\\sol, respectively, for wavelength-dependent radiation transfer and to 19.1~M\\sol, 20.1~M\\sol, and 22.9~M\\sol\\ for the corresponding simulations with grey radiation transfer. Our calculations demonstrate that massive stars can in principle be formed via accretion through a disk.  The accretion rate onto the central source increases rapidly after one initial free-fall time and decreases monotonically afterwards. By enhancing the non-isotropic character of the radiation field the accretion disk reduces the effects of radiative acceleration in the radial direction --- a process we denote the ``flashlight effect''.  The flashlight effect is further amplified in our case by including the effects of frequency dependent radiation transfer. We conclude with the warning that a careful treatment of radiation transfer is a mandatory requirement for realistic simulations of the formation of massive stars. ",
        "introduction": "Although massive stars play a critical role in the production of turbulent energy in the ISM, in the formation and destruction of molecular clouds, and ultimately in the dynamical and chemo-dynamical evolution of galaxies, our understanding of the sequence of events which leads to their formation is still rather limited. Because of their high luminosities we can expect: a) radiative acceleration will contribute significantly to the dynamical evolution during the formation process and b) the thermal evolution time scales of massive pre-main sequence objects will be extremely short. Thus, we cannot simply ``scale up'' theories of low mass star formation. Furthermore, OB stars form in clusters and associations; their mutual interactions via gravitational torques, powerful winds and ionizing radiation contribute further to the complexity of the problem.  Even though no massive disk has yet been directly observed around a main sequence massive star, it is likely that such disks are the natural consequence of the star formation process even in the high mass case. In their radio recombination maser studies and CO measurements Martin-Pintado et al. (1994) do find indirect evidence for both an ionized stellar wind and a neutral disk around MWC349. Moreover, several other high luminosity FIR sources --- suspected embedded young OB stars --- have powerful bipolar outflows associated with them (e.g., Eiroa et al. 1994; Shepherd et al. 2000).  Such massive outflows are probably powered by disk accretion, and, similar to their low mass counterparts, the flow energetics appear to scale with the luminosity of the source (see Cabrit \\& Bertout 1992; Shepherd \\& Churchwell 1996; Richer et al. 2000).  The detailed structure and evolutionary history of massive circumstellar disks has important consequences with regard to the early evolution of these protostars. Disks provide a reservoir of material with specific angular momentum too large to be directly accreted by the central object. Only after angular momentum is transported outwards can this material contribute to the final mass of the star. The transition region disk-star will strongly influence the star's photospheric appearance and how the star interacts with the disk. The relative high densities in these disks provide the environment for the further growth and evolution of dust grains, affecting the disk's opacity and consequently its energetics and appearance. The disk can be expected to interact with stellar outflows and is likely to be directly responsible for the outflows associated with star formation.  Disks surrounding massive stars or disks associated with close companions to massive stars should be short-lived compared to their low mass counterparts.  The UV environment within an OB cluster will lead to the photoevaporation of disks on a time scale of a few 10$^5$~yr (Hollenbach, Yorke, \\& Johnstone 2000).  Because this process operates on a time scale comparable to the formation of massive stars and is competitive to it, it is important to carefully model the transfer of radiation in the envelopes of accreting massive stars.  Numerical tools capable of this task are lacking at present.  We consider the present investigation as an important step in this direction. ",
        "conclusions": "Our improved frequency dependent radiation hydrodynamics code is able to track the infall of material within a molecular clump against radiative forces.  We find that the ``flashlight effect'' first discussed by Yorke \\& Bodenheimer (1999), i.e. the non-isotropic distribution of radiative flux that occurs when a circumstellar disk forms, is strongly compounded by the frequency dependent radiation transfer. The shortest wavelength radiation (which is also the most effective for radiative acceleration) is most strongly concentrated towards the polar directions, whereas the longer wavelength radiation (less effective radiative acceleration) is more or less isotropic. We conclude that massive stars can in principle be formed via accretion through a disk, in a manner analogous to the formation of lower mass stars.  A powerful radiation-driven outflow in the polar directions and a ``puffed-up'' (thick) disk result from the high luminosity of the central source.  We have developed a simplified model for following the evolution of accreting (proto-) stars, using existing tracks for non-accreting stars.  With this model we have shown that in the case of massive star formation the energy released within the accretion shock front, the ``accretion luminosity'', is {\\sl not} the dominant source of luminosity after a few thousand years of evolution.  The accretion rate onto the central source is time-dependent.  It rises sharply after one free-fall time to a maximum value and falls off gradually (in the frequency-dependent cases).  This is in contrast to the expectations of Meynet \\& Maeder (2000) and Behrend \\& Maeder (2001), who have assumed mass accretion rates $\\dot M_*$ that increase monotonically in time up to a maximum value.  We have also shown that the concept of ``birthline'', the equilibrium position of fully convective, deuterium-burning stars in the HR diagram with cosmic deuterium abundance, is --- strictly speaking --- unattainable for stars more massive than 1~M\\sol.  Beginning with a protostar of a fraction of a solar mass and building up via accretion to 1~M\\sol\\ and higher masses, it either accretes too rapidly (shifting the HR position to smaller radii) or it accretes too slowly (significant amounts of previously accreted deuterium are consumed).  For masses $M \\la 10$~M\\sol, however, the contribution of the accretion luminosity may make the star {\\sl appear} to lie on or above the birthline.  In this investigation we have not addressed the issues of the longevity of the circumstellar disk or the possible formation of a dense stellar cluster within our central computational zone rather than a single star.  However, even without the assumption of ionizing radiation, we find that these disks are not long-lived phenomena.  In the most massive cases the effects of radiative acceleration eventually disperse the remnant disks.  Future studies will have to address the issues of ionization and the interactions of the disk with powerful stellar winds.  The effects of nearby companions in a dense stellar cluster will also have to be considered in future work."
    },
    "0201/astro-ph0201331_arXiv.txt": {
        "abstract": "We model in simple terms the angular momentum ($J$) problem of galaxy formation in CDM, and identify the key elements of a scenario that can solve it.  The buildup of $J$ is modeled via dynamical friction and tidal stripping in mergers.  This reveals how over-cooling in incoming halos leads to transfer of $J$ from baryons to dark matter (DM), in conflict with observations.  By incorporating a simple recipe of supernova feedback, we match the observed $J$ distribution in disks.  Gas removal from small incoming halos, which make the low-$J$ component of the product, eliminates the low-$J$ baryons.  Partial heating and puffing-up of the gas in larger incoming halos, combined with tidal stripping, reduces the $J$ loss of baryons.  This implies a higher baryonic spin for lower mass halos.  The observed low baryonic fraction in dwarf galaxies is used to calibrate the characteristic velocity associated with supernova feedback, yielding $\\vfb \\sim 100\\kms$, within the range of theoretical expectations.  The model then reproduces the observed distribution of spin parameter among dwarf and bright galaxies, as well as the $J$ distribution inside these galaxies.  This suggests that the model captures the main features of a full scenario for resolving the spin crisis. ",
        "introduction": "\\label{sec:intro}  The `standard' model of cosmology, CDM, which assumes hierarchical buildup of structure, is facing difficulties in explaining observed properties of galaxies, such as the number density of dwarfs and the inner density profile of halos. Standing out is the angular-momentum problem, that is the apparent failure of the theory, via simulations, to reproduce the large sizes of disk galaxies and their structure. We make progress by first reproducing the problem via a simple model in which the important physical elements are spelled out, and then incorporating in this model the key process which may cure the problem --- feedback.  The sizes of disks are commonly linked to the spins of their parent halos as measured in N-body simulations [10]. The assumptions that the baryons and DM share the same distribution of specific angular momentum $j$ and that the baryons conserve their $j$ while contracting to a disk lead to disk sizes comparable to those observed.  However, simulations that incorporate gas find that most of the baryonic $j$ is transfered to the DM, resulting in disk sizes smaller by an order of magnitude [e.g.~14,15], and thus leading to a {\\it spin catastrophe}.  In addition, there is a {\\it mismatch of $j$ profiles}. The $j$ distribution (or profile) within simulated halos has been found to scatter about a universal shape, with an excess of low-$j$ (and high-$j$) material compared to the exponential disks observed [1, BD]. This mismatch has been demonstrated in an observed sample of 14 dwarf galaxies [18, BBS], which serves as the target for our modeling effort. BBS used for each halo the measured rotation curve and an assumed NFW profile to determine the halo virial quantities, with an average $\\langle \\vvir \\rangle \\simeq 60\\kms$. They then determined the baryonic spin parameter, averaging to $\\langle \\lpb \\rangle \\sim 0.07$, significantly larger than the $\\langle \\lpdm \\rangle \\sim 0.035$ of simulated halos, and then demonstrated the $j$-profile mismatch case by case.  BBS also estimated the ratio of disk to DM mass to be $\\langle \\fd \\rangle \\sim 0.04$, about a factor of 3 smaller than the universal fraction, indicating significant gas loss.  The spin catastrophe is commonly being associated with ``over-cooling\", that the gas rapidly cools and becomes tightly bound in small halos. When such a halo spirals into a bigger halo, the baryonic component survives intact all the way to the center and thus transfers all its orbital $j$ to the DM. It has therefore been speculated that energy feedback from supernova may remedy the problem by balancing the early cooling [e.g.~9]. A key idea is that the spin segregation between baryons and DM can go either way.  While gas cooling tends to lower the baryonic spin, heating due to feedback would reduce this effect, and gas removal from small halos would even lead to higher baryonic spin. However, a realistic implementation of feedback has proved challenging [e.g.~17]. The feedback process has not yet been studied or implemented in satisfactory detail. We do not know yet whether they can indeed solve the CDM problems, and how. This motivates our attempt to first understand how the feedback scenario may work using a very simple semi-analytic model. Knowing that in a hierarchical scenario the halo fromation can be largely interpreted as a sequence of mergers, our model is based on a simple algorithm for the buildup of halo spin by adding up the orbital angular momenta of merging satellites [13, MDS; 19]. It matches well the spin distribution among halos in N-body simulations as well as the $j$ profile within halos. This makes it a useful tool for understanding the over-cooling origin of the spin problem and for the attempt to cure it via feedback effects. Our work is described in more detail in [12, MD]. ",
        "conclusions": "\\label{sec:conc}  We devised a simple model to address the $j$ problems of galaxy formation within CDM.  By adding up the orbital ${\\bf J}$ in random realizations of merger histories, the model successfully reproduces the simulated distribution of spins among halos (MDS) and the distribution of $j$ within halos (MD).  A simple analysis of how the merger orbital $j$ turns into a spin profile provides a clue for how feedback effects in the satellite can resolve the spin problems. The idea is that the effective size of the gas component within the incoming halo determines its tidal stripping position in the big halo and thus its final remaining baryonic spin after the merger. The finding that the low-$j$ material originates in many minor mergers, that tend to cancel each other's ${\\bf J}$, provides the clue for a possible solution to the $j$-profile mismatch problem. The blowout of gas from small incoming halos, which is more pronounced in satellites of dwarfs, would eliminate the low-$j$ baryons in the merger product and increase the spin parameter, as observed.  The feedback effects, including heating and blowout, are modeled as a function of halo virial velocity, with one free parameter --- the characteristic velocity $\\vfb$ corresponding to the feedback energy from supernovae. To match the low baryonic fraction observed in dwarfs it has to be $\\vfb \\sim 100 \\kms$, consistent with the theoretical predictions [8]. This leads to an agreement between the model predictions and the observed disks, for the distribution of baryonic spin among galaxies and the baryonic $j$ distribution within galaxies, both dwarfs and bright galaxies.  We attempt to resolve the problems within the successful cosmological framework of CDM, by appealing to inevitable feedback effects. Another approach is to appeal to the Warm Dark Matter (WDM) scenario, despite the fact that it requires fine-tuning of the particle mass to $\\simeq 1~keV$.  The main feature of WDM is the suppression of small halos and the corresponding mergers.  While an N-body simulation of WDM [2] indicates the same $j$ properties of halos (the same properties can also be obtained as a general result of tidal-torque theory, see MDS), one expects the cooling to be less efficient in the absence of small halos, and thus the baryonic spin to be higher. However, the $j$ profile is still expected to be a problem, and the weaker feedback effects in the absence of small halos may not be enough for resolving it. These issues are yet to be studied in hydro simulations of WDM.  Feedback effects may also provide the cure to the missing dwarf problem in CDM, where the predicted number of dwarf halos is much larger than the observed number of dwarf galaxies [3]. While the number of dwarfs is automatically suppressed in WDM, it seems that the inclusion of the minimum inevitable feedback effects would reduce the predicted number of dwarfs to significantly below the observed number and thus be an overkill (J. Bullock, private comm.). Finally, we find [7] that the key elements of our toy model --- the tidal effects in mergers and the feedback in small halos --- are also very relevant in understanding and resolving the third problem of CDM, where the halos in simulations typically show steep cusps in their inner profiles [14], while observations indicate flat cores at least in some galaxies [4]. An analysis of tidal effects explains the inevitable formation of an asymptotic cusp as long as satellites continue penetrating into the halo center. Feedback effects may puff up small satellites, make them disrupt in the outer halo and thus allow a stable core.  The success of our toy model in matching several independent observations indicates that it indeed captures the relevant elements of the complex processes involved, and in particular that feedback effects may indeed provide the cure to all three problems of galaxy formation in CDM. The next natural step should be to incorporate a more sophisticated feedback recipe into the model using semi-analytic models and then full-scale cosmological simulations.    This research has been supported by the Israel Science Foundation grant 546/98, by the US-Israel Binational Science Foundation grant 98-00217, and by the German-Israeli Science Foundation grant I-629-62.14/1999."
    },
    "0201/astro-ph0201107_arXiv.txt": {
        "abstract": "{We present {\\em Chandra} observations of the galaxy cluster A4059.  We find strong evidence that the FR-I radio galaxy PKS~2354--35 at the center of A4059 is inflating cavities with radii $\\sim 20\\kpc$ in the intracluster medium (ICM), similar to the situation seen in Perseus A and Hydra A.  We also find evidence for interaction between the ICM and PKS~2354--35 on small scales in the very center of the cluster.  Arguments are presented suggesting that this radio galaxy has faded significantly in radio power (possibly from an FR-II state) over the past $10^8\\yr$.} ",
        "introduction": "Clusters of galaxies are complex dynamical structures and their cores are subject to an array of interesting physical processes.  Constraints from imaging X-ray observations suggest that the hot X-ray emitting intracluster medium (ICM) in the core regions of rich clusters is radiatively cooling on timescales shorter than the life of the cluster, giving rise to cooling flows \\citep[][and references therein]{fabian:94}.  The central dominant galaxy present in many clusters often hosts a radio loud active galactic nucleus (AGN).  It has been suggested \\citep[e.g.,][]{binney:95} that cooling flows and central cluster radio galaxies are intimately related via complex feedback processes.  It is easy to see how radio galaxy activity resulting from black hole accretion can be associated with a cooling flow. However, the impact of a radio galaxy on its environment is much less clear.  Theoretically, we expect radio jets to inflate cocoons of relativistic plasma that expand into the surrounding ICM \\citep[e.g.][, hereafter RHB]{begelman:89,kaiser:97,reynolds:01}.  The energy input by this process has recently come under investigation for its potential role in heating cluster cores \\citep{bruggen:01,quilis:01,reynolds:01b}.  However, while our simulations suggest that about half of the energy injected by the jets can be thermalized in the cluster center, numerical simulations of this process still carry a large degree of uncertainty, since limited computational resources require significant simplications.  In order to verify the validity of the assumptions and to design future models, we require guidance from observations of radio-galaxy/cluster interactions.  Imaging X-ray observatories, such as the {\\it Chandra X-ray Observatory} (CXO), provide a direct probe of this interaction.  Both {\\it ROSAT} and CXO observations of Perseus~A have found X-ray cavities coincident with the radio lobes \\citep{boehringer:93,fabian:00}, surrounded by X-ray shells which appear to be slightly cooler than the unperturbed ICM (see, for example, RHB for a possible explanation).  Similar features are seen in CXO observations of Hydra~A \\citep{mcnamara:00,david:01} and Abell 2052 \\citep{blanton:01}.  In this {\\it Letter}, we present CXO observations of the rich galaxy cluster Abell~4059 ($z=0.049$).  The cD galaxy of A4059 hosts the FR-I radio galaxy PKS~2354--35.  A short {\\it ROSAT} High Resolution Imager (HRI) observation of this source suggested the presence of two ICM cavities at the same position angle as the radio lobes \\citep[][, hereafter HS]{huang:98}.  In \\S~2 we discuss our observations, confirming the presence of these cavities, and show that A4059 displays significant additional morphological complexity.  Constraints on models for this source are discussed in \\S~3, \\S~4 presents our conclusions.  We assume a Hubble constant of $H_0=65 \\kmpspMpc$ and $q_0 = 1/2$, giving a linear scale of $1\\, {\\rm kpc\\,arcsec^{-1}} = 0.492\\, {\\rm kpc\\,pixel^{-1}}$. ",
        "conclusions": "We have presented CXO observations of Abell~4059.  While the ICM appears smooth and relaxed on large scales, it shows complex morphology in the core region which is likely the result of interaction between the ICM and the central FR-I radio galaxy PKS2354--35.  As was suggested by HS, PKS~2354--35 appears to have inflated two large cavities in the ICM. Together with a central bar-like structure, these cavities produce an hour-glass like morphology which can be readily understood as being due to a radio cocoon expanding into the ICM.  While clear correspondence exists between the NW cavity and the NW radio lobe, the SE cavity is much larger than the SE lobe, suggesting that this could be a `missing link' between cavities with and without visible radio lobes.  The absence of sharp edges in the brightness images and of large temperature jumps implies that PKS2354--35 is {\\it not} driving a strong shock into the ICM.  We suggest that it is in the weak-shock/compression-wave phase identified in the hydrodynamic simulations of RHB.  Dynamical estimates give a time averaged kinetic source power of at least $L_{\\rm kin} \\gtrsim 3\\times 10^{44}\\ergps$, while estimates based on the current radio luminosity indicate a source power of $L_{\\rm kin} \\lesssim 7\\times 10^{43}\\ergps$.  We suggest that this source has faded by a significant amount (and possibly from an FR-II phase) during the past $10^8\\yr$."
    },
    "0201/astro-ph0201382_arXiv.txt": {
        "abstract": "We present spectroscopic and imaging observations of the active T dwarf 2MASS 1237+6526, intended to investigate the emission mechanism of this cool brown dwarf. The H$\\alpha$ emission line first detected in 1999 July appears to be persistent over 1.6 years, with no significant variation from $\\log_{10}$(L$_{H{\\alpha}}$/L$_{bol}$) = $-$4.3, ruling out flaring as a possible source. The relatively high level of emission in this object appears to be unique amongst observed late-L and T dwarfs. One of our spectra shows an apparent velocity shift in the H$\\alpha$ line, which could be indicative of an accretion hot spot in orbit around the brown dwarf; further confirmation of this shift is required. J-band monitoring observations fail to detect any significant variability (e.g., eclipsing events) at the $\\pm$0.025 mag level over periods of up to 2.5 hours, and there appears to be no statistical evidence of variability for periods of up to 14 hours. These limits constrain the mass of a hypothetical interacting secondary to M$_2$ $\\lesssim$ 20 M$_{Jup}$ for inclinations $i$ $\\gtrsim$ 60$\\degr$. While our observations do not explicitly rule out the binary hypothesis for this object, it does suggest that other mechanisms, such as youthful accretion, may be responsible. ",
        "introduction": "A previously undetected population of T dwarfs, brown dwarfs exhibiting CH$_4$ absorption bands at 1.6 and 2.2 $\\micron$ \\citep{kir99,me01a}, have recently been identified by various deep optical and near-infrared surveys \\citep{str99,me99,me00a,me00c,me01a,cub99,tsv00,leg00,geb01b}.  One of these objects, 2MASSI J1237392+652615 \\citep[hereafter, 2MASS 1237+6526]{me99}, identified in the Two Micron All Sky Survey \\citep[hereafter 2MASS]{skr97}, was found to have H$\\alpha$ in emission, with a relative luminosity $\\log_{10}$(L$_{H{\\alpha}}$/L$_{bol}$) = $-$4.3 \\citep{me00b}. While many low-mass stars and young brown dwarfs are seen to exhibit this feature, likely generated by chromospheric magnetic activity, the frequency and relative luminosity of H$\\alpha$ emission drops rapidly beyond spectral types M7 V \\citep{kir00,giz00}. Indeed, at the time of its detection, the emission of 2MASS 1237+6526 was unique amongst objects later than type L5 V. \\citet{kir01} have recently reported weak H$\\alpha$ emission in the bright T2 V SDSSp J125453.90-012247.4 \\citep[hereafter SDSS 1254-0122]{leg00}.  In order to explain the activity of 2MASS 1237+6526, \\citet{me00b} have hypothesized that this object may actually be an interacting brown dwarf binary system.  In this scenario, the binary mass ratio and orbital separation are such that the lower-mass secondary fills its Roche lobe and steadily loses mass to the primary. For a 70 M$_{Jup}$\\footnote{We adopt the definitions M$_{Jup}$ = 1.90$\\times$10$^{30}$ g = 9.55$\\times$10$^{-4}$ M$_{\\sun}$, R$_{Jup}$ = 7.15$\\times$10$^{9}$ cm = 0.10 R$_{\\sun}$ \\citep{cox00,tho00}.} primary and mass ratio $q$ $\\equiv$ $M_2/M_1$ $>$ 0.07 (i.e., $M_2$ $>$ 5 M$_{Jup}$), sustained mass loss can occur for physical separations $a$ $\\lesssim$ 6 R$_{Jup}$ and orbital periods $p$ $\\lesssim$ 5.5 hours.  The close orbit required for this scenario currently rules out direct imaging of the binary pair ($a$ $<$ 0$\\farcs$0002 assuming a distance of 14 pc; Burgasser et al.\\ 1999); however, partial eclipsing could be observed in a system with an inclination $i$ $\\gtrsim$ 60$\\degr$.  We have obtained followup spectroscopic and photometric observations of 2MASS 1237+6526 in order to test the binary hypothesis, as well as examine the possibility that this object was originally observed during a prolonged flare. In $\\S$2, we present red optical (6300--10100 {\\AA}) spectral data obtained over 1.6 yr using the Keck 10m Low Resolution Imaging Spectrograph \\citep[hereafter LRIS]{oke95}, and discuss the behavior of the H$\\alpha$ emission line over this period. J-band monitoring observations obtained over two nights using the Palomar 60'' Near-Infrared Camera \\citep[hereafter IRCam]{mur95} are presented in $\\S$3. We discuss how the imaging observations constrain the binary scenario in $\\S$4.  Finally, in $\\S$5 we discuss how these follow-up observations constrain various emission mechanisms for this unique object. ",
        "conclusions": "Based on the results of the monitoring observations, we can confidently rule out an interacting binary system for 2MASS 1237+6526 for $i$ $>$ 60$\\degr$ and M$_2$ $>$ 20 M$_{Jup}$.  We cannot rule out a less inclined system, although such a geometry would fail to explain the apparent line shift seen in our 2000 March spectral data. A compromise solution would be a moderately inclined system with a massive primary and large mass ratio.  Conclusive synthesis of the spectroscopic and photometric data, however, requires verification of H$\\alpha$ line shift.  Other emission mechanisms should also be considered. 2MASS 1237+6526 could be a young, and hence very low mass, brown dwarf that is still accreting material from a circum(sub)stellar disk. If this object were in a weak-lined T-Tauri (WLTT) phase, its H$\\alpha$ line could arise from a small, optically thin accretion boundary layer, consistent with our spectral and photometric observations. Assuming a typical WLTT age (30--100 Myr; Hartmann, Kenyon, \\& Hartigan 1993) and T$_{eff}$ $\\sim$ 950 K, the evolutionary models of \\citet{bur97} predict a mass of 3--12 M$_{Jup}$, below the Deuterium-burning limit.  Note that 2MASS 1237+6526 is not associated with any known star-forming region, although it could have been ejected from its nascent cluster \\citep{rep01}.  The absence of the 1.25 $\\micron$ K I doublet in this object (I.\\ McLean, priv.\\ comm.) marginally supports the possibility of it being a low-gravity source, although we cannot rule out a purely temperature effect \\citep{me01a}.  Our follow-up observations of 2MASS 1237+6526 have enabled us to rule out flaring as a viable means of emission, and have placed significant constraints on the geometry and membership of an interacting binary system. Parallax (to determine absolute brightness) and space motion measurements are clearly required to further characterize this object. Furthermore, the apparent velocity shift described in $\\S$2.4, which may prove to be a vital clue to the origin of emission in 2MASS 1237+6526, requires confirmation and more detailed, higher-resolution investigation. The nature of this enigmatic brown dwarf may hopefully be revealed with further observations."
    },
    "0201/astro-ph0201457_arXiv.txt": {
        "abstract": "There is much evidence to suggest that stellar wind capture, rather than Roche lobe overflow, serves as the accretion mechanism onto the compact secondary object in the massive X-ray binary LS~5039.  The lack of significant emission combined with only a modest X-ray flux provide observational evidence that no large-scale mass transfer is occurring (consistent with our estimate of the radius of the O6.5~V((f)) optical star that is smaller than its critical Roche radius).  Here we determine the mass loss rate of the optical star from the broad, residual emission in the H$\\alpha$ profile. Using a stellar wind accretion model for a range in assumed primary mass, we compute the predicted X-ray luminosity for the system.  We compare our results to the observed X-ray luminosity to determine the mass of the compact object for each case.  The companion appears to be a neutron star with a mass between 1 and $3 M_\\odot$.  With our new constraints on the masses of both components, we discuss their implications on the evolution of the system before and after the supernova event that created the compact companion.  The binary experienced significant mass loss during the supernova, and we find that the predictions for the resulting runaway velocity agree well with the observed peculiar space velocity. LS~5039 may be the fastest runaway object among known massive X-ray binaries. ",
        "introduction": "%  LS~5039 is a relatively faint ($V = 11.3$) and massive star of type O6.5~V((f)) \\citep{cla01} that is one of only a few confirmed massive X-ray binaries (MXRBs) with associated radio emission \\citep{rib99}. It has radio-emitting relativistic jets characteristic of galactic microquasars, and it is probably a high energy gamma ray source as well \\citep{par00}. We recently discovered that the system is a short period binary ($P = 4.117 \\pm 0.011$ d) with the highest known eccentricity ($e = 0.41\\pm0.05$) among O star binaries with comparable periods \\citep{mcs01a}. This high eccentricity probably results from the huge mass loss that occurred with the supernova (SN) explosion that gave birth to the compact star in the system \\citep{bha91,nel99}. Binaries that suffer large mass loss in a SN are expected to become runaway stars, and recently both we \\citep{mcs01b} and \\citet{rib02} found that LS~5039 has a relatively large proper motion that indicates that the binary has a record-breaking peculiar space velocity among MXRBs.  Reliable estimates for the masses of the components of the binary are of key importance for any discussion about the evolution of this remarkable system.  Here we present an investigation of the possible mass range for the X-ray star based on a stellar wind accretion model for the X-ray production (\\S2).  This analysis relies on our observations of the wind emission effects in the H$\\alpha$ profile and the wind models of \\citet{pul96}.   We then apply our derived mass estimates for both stars to determine the probable masses and orbital parameters prior to the SN based on the system's current eccentricity (\\S3).   We find that the predicted and observed runaway velocities are in good agreement, and, thus, LS~5039 provides the best verification to date of model predictions about the outcome of a SN explosion in a massive binary. ",
        "conclusions": ""
    },
    "0201/astro-ph0201219_arXiv.txt": {
        "abstract": "We have used the WFPC2 camera of the {\\sl Hubble Space Telescope\\/} to obtain deep F814W images of a blank field in the Virgo Cluster located $41\\arcmin$ northwest of M87.  We perform star counts in that field, and in another Virgo field observed by \\citet{ftv98}, and show that, when compared to the Hubble Deep Field North and South, the Virgo Cluster contains an excess of objects with magnitudes $I \\gtrsim 27$.  We attribute this excess to a population of intracluster red-giant branch (IC-RGB) stars.  By modeling the luminosity function of these stars, we show that the tip of the Virgo RGB is at $I_{TRGB} \\sim 27.31^{+0.27}_{-0.17}$ and that the cluster contains a small, but significant, excess of stars that are up to $\\sim 1$~mag brighter than this tip.  If this luminous component is due entirely to stars on the asymptotic giant branch (AGB), it implies an age for the population of $> 2$~Gyr; if foreground RGB stars contribute to the luminous tail, then the derived age for the stars is older still.  The luminosity function also suggests that most of the intracluster stars are moderately metal-rich ($-0.8 \\lesssim$ [Fe/H] $\\lesssim -0.2$), a result consistent with that expected from stars that have been tidally stripped from intermediate luminosity galaxies.  Additionally, a comparison with the planetary nebulae in our field also supports this view, although the existence of a more metal-poor population (from stripped dwarfs) cannot be ruled out. Our derived average surface brightness, $\\mu_I = 27.9^{+0.3}_{-0.5}$ mag arcsec$^{-2}$ for Virgo's diffuse component suggests that intracluster stars contribute 10\\% to 20\\% of the cluster's total $I$-band luminosity. ",
        "introduction": "Intracluster stars (stars associated with galaxy cluster {\\it potentials,} rather than with any particular galaxy) provide an important clue towards the understanding of the formation and evolution of galaxies and galaxy clusters. N-body simulations show that this diffuse component can be produced by a number of processes.  For instance, tidal interactions between merging galaxies \\citep{mil83,weil97,dub99}, between a galaxy and the cluster potential \\citep{mer84,dub99}, and between galaxies during high speed encounters \\citep[\\ie `galaxy harassment';][]{moore96, moore98} can all liberate stars \\citep[and globular clusters, \\eg][]{west95} into intracluster space. Alternatively, a significant number of intracluster stars may be created early on during a cluster's initial collapse \\citep{mer84}.  By observing these stars and determining their photometric and kinematic properties, we can therefore learn about the workings of tidal-stripping, the distribution of dark matter around galaxies, and the initial conditions of cluster formation.  Unfortunately, observational studies of intracluster light are very difficult due to its very low surface brightnesses (typically $\\mu_B \\gtrsim 27$~mag arcsec$^{-2}$, or less than 1\\% of the background sky).  Consequently, though the first detection of diffuse intracluster light was made a full half-century ago \\citep{zw51} and there have been numerous studies thereafter \\citep[\\eg][]{oem73,mat77,mel77,tk77,uson91,vg94,bern95,gonz00}, there is little agreement about the most basic data.  For example, even in the well-observed Coma Cluster, measurements of the fraction of intracluster light range from less than 25\\% \\citep{mel77} to $\\sim 50\\%$ \\citep{bern95} of the total cluster luminosity.  An obvious complement to measurements of the diffuse light in clusters is the direct detection and measurement of individual intracluster stars. Although this is only possible in nearby clusters (\\eg Virgo, Fornax, Centaurus), investigations of individual stars have the advantage of removing many sources of error that typically complicate surface brightness measurements (\\ie contamination by low-surface brightness dwarf galaxies, scattered light from foreground stars, flat-fielding errors, etc.).  Moreover, by studying individual stars, one has the hope of determining the underlying population's age, metallicity, and dynamical properties.  Because of their probative value, searches for intracluster stars have become common in recent years.  In particular, a number of wide-field on-band/off-band [O~III] $\\lambda 5007$ imaging surveys for intracluster planetary nebulae (IPN) have been conducted in fields of the Virgo and Fornax Clusters \\citep{tw97, men97, feld98, feld00}.  These studies have confirmed the existence of large numbers of intracluster stars, and have produced evidence to suggest that many of these stars are of moderate age and metallicity.  Although IPN are a powerful probe of intracluster starlight, they do have some limitations.   Spectroscopy of Virgo IPN candidates by \\citet{kud00} and \\citet{free00}, as well as a blank-field imaging survey by \\citet{rbc01} have shown that not all objects detected through narrow-band $\\lambda 5007$ filters are planetary nebulae -- about 20\\% of the detections in Virgo appear to be Ly$\\alpha$ galaxies at $z = 3.13$.  This source of contamination produces an ambiguity in the IPN analysis, which can only be broken via time-consuming spectroscopy.  In addition, in order to determine the total amount of intracluster light from IPN observations, one needs to know the production rate of bright planetaries normalized to the bolometric luminosity of the stellar population.  Observations demonstrate that this quantity varies by almost an order of magnitude depending on the stellar population \\citep{peim90,rbc95}; this creates a fundamental problem of the interpretation the IPN observations. Finally, planetary nebulae are relatively rare objects: typically it takes $\\sim 5 \\times 10^8 L_\\odot$ of stars to produce $\\sim 1$ [O~III] bright planetary.  An alternative approach to studying intracluster starlight is to search for the constituent red giant (RGB) and asymptotic giant branch (AGB) stars of the stellar population.  RGB stars are much more numerous than planetary nebulae, and therefore surveys for their presence do not require wide-field telescopes. Moreover, translating the number counts of red giants to total population luminosity is much more straightforward: there is no ambiguity as with PN production.  Finally, because the absolute magnitude of the red giant branch tip is a function of metallicity (for relatively high metallicity populations), the luminosity function of RGB and AGB stars allows us to constrain both the metallicity and age of the stellar population.  \\citet{ftv98} (hereafter FTV) were the first to detect individual RGB stars in intracluster space.  By using the WFPC2 camera of the {\\sl Hubble Space Telescope\\/} to take deep F814W ($I$-band) images of a ``blank'' field located 45$^\\prime$ E of the central Virgo cluster galaxy M87, Ferguson \\etal were able to detect the presence of intracluster red giants through the statistical excess of point sources over that seen in the Hubble Deep Field North. From their data, Ferguson \\etal concluded that intracluster stars make up $\\sim 10\\%$ of the total stellar mass of the system.  Unfortunately, the Ferguson \\etal analysis was necessarily limited.  Because their survey area consisted of only one WFPC2 field, Ferguson \\etal found only a small number of stars, and thus could not place significant constraints on the age or metallicity of the intracluster population.  Moreover, with only the one field, Ferguson \\etal could not address the question of the overall distribution of these stars.  Recent discoveries of low-surface brightness arcs in other nearby clusters \\citep{gw98,tm98,cr00} and large field-to-field variations in the number density of Virgo IPN \\citep{feld98} have demonstrated that intracluster stars are not distributed uniformly. Consequently, to constrain the underlying population of these stars and learn about their large-scale distribution, additional fields must be studied.  Here we present the results of a second study of Virgo intracluster RGB stars. We begin by describing our deep {\\sl HST\\/} observations of a Virgo blank field located between M87 and M86, near the center of the cluster's ``sub-clump A'' \\citep{bts87}.  We detail our reduction techniques, our photometric procedures, and the artificial star simulations needed to measure the errors and incompleteness of our measurements.  In section 3, we combine our data with those of the FTV survey, and compare the raw Virgo Cluster stellar luminosity function (LF) with that of the North and South Hubble Deep Fields.  We show that Virgo possesses a significant excess of point sources that is due to the cluster's population of RGB and (possibly) AGB stars.  In section 4, we model the point-source luminosity function, and place constraints on the cluster's distance, and on the age and metallicity of its intracluster population(s). Finally, we discuss these results and compare them with other measures of intracluster stars. ",
        "conclusions": "We have analyzed deep F814W {\\sl HST\\/} images of a single Virgo cluster field located 41$^\\prime$ NW of M87, near the cluster center.   Photometry of the unresolved objects in this field \\citep[combined with data from another Virgo cluster field observed by][]{ftv98} shows an excess of objects (with respect to the background HDF-N and HDF-S fields) with $I \\gtrsim 27$, which we attribute to intracluster RGB stars in the Virgo cluster.   We derive an average surface brightness of $\\mu_I = 27.9^{+0.3}_{-0.5}$ mag~arcsec$^{-2}$ for both fields; if our data are representative of the cluster's IC light in general, then IC stars comprise $15^{+7}_{-5}\\%$ of Virgo's total light.   This result is similar to that obtained from observations of IC planetary nebulae for values of $\\alpha_{2.5} = 23^{+10}_{-12} \\times 10^{-9}$ PN $L_{\\odot}^{-1}$.  We have modelled the resulting luminosity function with a single-component RGB+AGB population, and derived the location of both the RGB tip ($I_{TRGB}= 27.31^{+0.27}_{-0.17}$) and the bright extent of an AGB ($\\Delta I = 0.8^{+0.2}_{-0.2}$).  We note, however, that the latter is probably contaminated by foreground RGB stars.  We find that the RGB tip is significantly fainter than that observed in a Virgo cluster dE,N galaxy \\citep{har98}, and suggest that this difference is due to a higher metal abundance for the intracluster stars ($-0.8\\lesssim$ [Fe/H] $\\lesssim -0.2$). Our measurement of the intracluster AGB population indicates that the stars are old ($t > 2$~Gyr), but due to the possible existence of a foreground RGB component, we cannot place a firm limit on the population age.  From our observations, it seems most likely that the bulk of Virgo's intracluster stars were once stripped from lower-mass spiral and elliptical galaxies, but we cannot rule out the possibility that a significant metal-poor population (such as that expected from tidally stripped {\\it dwarf\\/} galaxies) exists. It is clear that measurements of the metallicity distribution of IC stars will be the key to understanding their origins."
    },
    "0201/astro-ph0201443_arXiv.txt": {
        "abstract": "The theoretical uncertainties in the calibration of the relationship between the subgiant mass and age in metal-poor stars are investigated using a Monte Carlo approach.  Assuming that the mass and iron abundance of a subgiant star are known exactly, uncertainties in the input physics used to construct stellar evolution models and isochrones lead to a Gaussian 1-$\\sigma$ uncertainty of $\\pm 2.9\\%$ in the derived ages.  The theoretical error budget is dominated by the uncertainties in the calculated opacities.  Observations of detached double lined eclipsing binary OGLEGC-17 in the globular cluster \\wcen\\ have found that the primary is on the subgiant branch with a mass of $\\mathrm{M} = 0.809\\pm 0.012\\,\\mathrm{M}_{\\odot}$ and $\\feh = -2.29\\pm 0.15$ \\citep{kaluzny}. Combining the theoretical uncertainties with the observational errors leads to an age for OGLEGC-17 of $11.10\\pm 0.67\\,$Gyr.  The one-sided, 95\\% lower limit to the age of OGLEGC-17 is 10.06 Gyr, while the one-sided, 95\\% upper limit is 12.27 Gyr. ",
        "introduction": "Traditionally, absolute globular cluster (GC) ages have been determined using the absolute magnitude of the main sequence turn-off (TO), or subgiant branch (SGB), as this minimizes the theoretical uncertainties associated with stellar evolution models \\citep[e.g.\\ ][]{renz91,mvsgb}.  This age determination method requires that the distance to the GC be known.  There is considerable uncertainty regarding the distance scale to GCs, and this translates into a significant uncertainty in the absolute age estimates of GC \\citep{mc3}. To avoid this error \\cite{pacz} has advocated the use of detached eclipsing double line spectroscopic binaries to determine the age of GCs.  In these binary systems, it is possible to determine the mass of the individual stars.  These mass estimates are derived in a fundamental manner, and are likely to be free from systematic errors \\citep{pacz}. If one of the members of the binary is at the TO, or on the SGB then the age of the cluster may be determined from the TO/SGB mass-age relation.  In principle the relation between the TO/SGB mass and age is robust prediction of stellar evolution theory -- it simply depends on the amount of hydrogen fuel available for nuclear burning in the core of the star and the luminosity of the star during its main sequence lifetime.  Thus, the TO/SGB mass-age relation should be insensitive to the details of what occurs near the surface of stars and will not depend on the treatment of convection for the low mass stars in GCs \\citep{pacz}.  For these reasons, one might expect that ages derived from the masses of TO/SGB stars will be relatively insensitive to various significant uncertainties that might otherwise be important in stellar structure calculations.  This paper will explore how the uncertainties in stellar structure and evolution calculations (\\S \\ref{uncertain}) translate into errors in ages derived from SGB masses in GCs (\\S \\ref{monte}).  This work is motivated by the high precision mass estimate for the detached eclipsing double line spectroscopic binary OGLEGC-17 in \\wcen\\ by \\cite{kaluzny}. The primary in OGLEGC-17 is on SGB \\citep{thompson}. The age of this star is derived in \\S \\ref{age}, and this paper concludes with a general discussion of the implications of this age determination in \\S \\ref{universe}. ",
        "conclusions": "} The age of OGLEGC-17 may be compared to our  estimate of the mean age of 17 metal-poor GCs which used the luminosity of the TO as an age indicator \\citep{mc3}. For the same set of input parameters, we found a median age of 12.5 Gyr, and one sided 95\\% confidence level ages of 10.2 Gyr and 15.9 Gyr.  The non-Gaussian distribution has a  lower $1\\,\\sigma$  age of 11.0 Gyr, implying that age of OGLEGC-17 and the mean age of 17 metal-poor GCs agree at the $1\\,\\sigma$ level.  The one-sided 95\\% confidence level lower limits to the two age determinations are quite similar (10.1 and 10.2 Gyr).  This supports our conclusion that the ages of the oldest stars and recent measurements of the Hubble constant require that the cosmic equation of state has $w \\equiv \\textrm{pressure/density} < -0.3$ \\citep{mc3}.  The age of OGLEGC-17 was determined assuming that the error in the mass determination was Gaussian.  As discussion of the error in the mass determination of OGLEGC-17 has not been published it is not clear if this assumption is valid.  If it is, then age of OGLEGC-17 is known much more accurately than the mean age of the metal-poor GCs determined from their TO luminosity.  The upper limit on the mean age (12.3 Gyr) is much smaller than that determined in the GC study (15.9 Gyr).  The upper limit to the age of OGLEGC-17 may be compared to the age of the universe determined from the cosmic microwave background of $14.0\\pm 0.5\\,$Gyr \\citep{knox}.  Their $2\\,\\sigma$ lower limit of $13.0$ Gyr is 0.7 Gyr older than our upper limit, implying at least 0.7 Gyr of galaxy evolution before OGLEGC-17 formed.  This corresponds to a redshift of globular cluster formation of $z \\la 7$ (cf.\\ equation 1 in \\cite{mc3}).  It is worth remarking that when more old GC ages are constrained in this way, a comparison strict upper limits one might derive on their ages with the Hubble age may provide the strongest constraints on cosmological models with exotic forms of dark energy such that $w= <-1$."
    },
    "0201/astro-ph0201169_arXiv.txt": {
        "abstract": "This paper presents ASCA/SIS and ROSAT/HRI results of three supernova remnants (SNRs) in the Small Magellanic Cloud: 0103$-$726, 0045$-$734, and 0057$-$7226. The ROSAT/HRI images of these SNRs indicate that the most of the X-ray emissions are concentrated in the center region. Only from 0103$-$726 are faint X-rays along the radio shell also detected. The ASCA/SIS spectra of 0103$-$726 and 0045$-$734 exhibit strong emission lines from highly ionized metals. The spectra were well-fitted with non-equilibrium ionization (NEI) plasma models. The metal abundances are found to be larger than the mean chemical compositions in the interstellar medium (ISM) of the SMC. Thus, X-rays from these two SNRs are attributable to the ejecta gas, although the ages estimated from the ionization timescale  are significantly large, $\\gtrsim 10^4$~yr. The chemical compositions are roughly consistent with the type-II supernova origin of a progenitor mass $\\lesssim 20 M_{\\solar}$. The SIS spectrum of 0057$-$7226 was also fitted with an NEI model of an estimated age $\\gtrsim 6 \\times 10^3$~yr. Although no constraint on the metal abundances was obtained, the rather weak emission lines are  consistent with the low metal abundances in the ISM of the SMC. A possible scenario for the evolution  of the morphologies and spectra of SNRs is proposed. ",
        "introduction": "\\label{sec:intro}  The Large and Small Magellanic Clouds (hereafter LMC and SMC) are ideal galaxies  for systematic X-ray studies of SNRs, because of their reasonable sizes, well-calibrated distances (50~kpc and 60~kpc, respectively; \\cite{ber2000}), and low interstellar absorptions.  \\citet{wil1999} obtained X-ray images of LMC SNRs mainly with ROSAT/HRI, and classified the SNRs into six classes according to the morphology. They found a loose correlation between the SNR sizes and the morphological classes: the smaller SNRs tend to exhibit shell-like X-ray emission, while the larger ones have ``diffuse-face'' or centrally-brightened morphologies. This correlation indicates the evolution from shell-like to centrally-brightened SNRs. \\citet{wil1999} argued that the centrally brightened morphologies could be produced by the cloud evaporation model \\citep{whi1991} or the fossil radiation model  (\\cite{rho1998}, and references therein).  \\citet{hug1998} systematically studied the ASCA/SIS spectra of seven middle-aged SNRs with a model assuming an internal structure based on the Sedov solution. They found that the average metal abundances in the SNRs decrease with increasing SNR age, and gradually approach to the mean value of the interstellar medium (ISM) in the LMC. \\citet{nis2001} analyzed nine fainter LMC SNRs in a less-systematic way and determined the basic plasma parameters. Besides the general tendency of the abundance decrease, five SNRs among the 16 SNRs examined by \\citet{hug1998} and \\citet{nis2001} show overabundances of several elements (overabundant species are different for each SNR).   In the SMC, on the other hand, detailed X-ray studies have been limited to the brightest and youngest SNR, 0102$-$723, due mainly to the far-fainter X-rays from the other SNRs. \\citet{hay1994} analyzed the ASCA/SIS spectrum from the full region of 0102$-$723 and found overabundances of heavy elements, consistent with a type-II young SNR. With the Chandra ACIS, \\citet{hug2000} have spatially resolved X-rays from the blast-wave and those from the ejecta-dominated inner region; the former exhibits low elemental abundances consistent with the SMC ISM, while the latter is overabundant, especially in O and Ne.   The second-brightest X-ray SNR in the SMC is 0103$-$726, and the next luminous class includes several SNRs, such as  0045$-$734 and 0057$-$7226 \\citep{hab2000}. All of these SNRs are located in H~{\\sc ii} regions, DEM~S125, N19, and N66, respectively (e.g., \\cite{fil1998}). \\citet{mil1982} carried out radio observations with the Molonglo Observatory synthesis telescope (MOST) and determined the diameters of the former two SNRs to be $\\sim 52$~pc (0103$-$726) and $\\sim 25$~pc (0045$-$734). The Einstein/HRI image of 0103$-$726 shows extended X-rays with the same size of radio emission \\citep{ino1983}.  In order to distinguish SNRs from the strong background emission of the H~{\\sc ii} region, \\citet{ye1991} made a map of the intensity difference between H$\\alpha$ and the radio continuum, instead of the conventional method of using the intensity ratio of [S~{\\sc ii}]/H$\\alpha$. They then found a non-thermal shell structure with a diameter of  $\\sim 56$~pc, the SNR 0057$-$7226, from the most luminous H~{\\sc ii} region in the SMC, N66. The sizes of these three SNRs are thus much larger than that of 0102$-$723, $\\sim 8$~pc \\citep{mil1982}. By an analogy of the LMC samples used by \\citet{hug1998}, these large sizes indicate that the three SNRs should be very old, possibly older than $10^4$~yr, or even more. In fact, \\citet{ros1994} carried out a  kinematic study in H$\\alpha$ emission and determined an age of 0045$-$734 to be $5.7 \\times 10^4$~yr with a Sedov phase assumption. Usually, X-ray spectra from such large (old) SNRs are expected to exhibit low metal abundances consistent with the ISM, as has already been reported in the  LMC samples \\citep{hug1998}.   However, a preliminary analysis of the ASCA spectra by \\citet{yok2000} does not agree with the above prediction, except for 0057$-$7226 (N66): overabundances of several elements were suggested for 0103$-$726 and 0045$-$734 (N19). We therefore carried out more detailed and elaborate X-ray studies on the three SNRs with a uniform analysis, using the high-resolution spectroscopy and imaging data of the ASCA/SIS and ROSAT/HRI, as well as information from radio observations. ",
        "conclusions": "\\label{sec:dis} \\subsection{Plasma Parameters in the SNRs}\\label{subsec:plasma}  Using the X-ray images (figure \\ref{fig:image}), we estimated the volume $V$ of the X-ray emitting plasma. For simplicity, we assumed that the plasma in 0103$-$726 and 0057$-$7226 fills a sphere with a diameter of \\timeform{2.5'} and $2'$, respectively, while that in 0045$-$734 fills a cylinder with a height of $3'$ and a diameter of $1'$. We then obtained $V =  1.3 \\times 10^{60} \\cdot \\beta$~cm$^{3}$ for 0103$-$726, $V =  3.6 \\times 10^{59} \\cdot \\beta$~cm$^{3}$ for 0045$-$734, and $V =  6.6 \\times 10^{59} \\cdot \\beta$~cm$^{3}$ for 0057$-$7226, respectively, where $\\beta$ represents the volume filling factor, which could be different for each SNR.   From  the emission measure $EM$ determined by the spectral fitting, we could estimate the electron density of the plasma, $n_{\\rm e}$, by $n_{\\rm e} =  \\sqrt{EM/V}$. The number density of the nucleons was simply assumed to be the same as that of electrons. The age $t$ was then determined from the ionization timescale, $\\tau$, by $t =  \\tau/n_{\\rm e}$. The total mass of the plasma $M_{\\rm total}$ was estimated by  $M_{\\rm total} =  n_{\\rm e}Vm_{\\rm H}$, where $m_{\\rm H}$ is the mass of a hydrogen atom. We then determined the mass of the overabundant heavy elements, using $M_{\\rm total}$ and the best-fit abundances. These estimated parameters are given in table \\ref{tab:plasma}. The lower limit for the ionization age of 0045$-$734 is  $2.9\\times 10^4$~yr in both models, which is consistent with the age of $5.7 \\times 10^4$~yr derived from an optical study of the kinematics \\citep{ros1994}.  For 0103$-$726 and 0045$-$734, the abundances are significantly larger than the SMC mean value in any model; hence, the ejecta should largely contribute to the X-ray emitting plasma. In figure \\ref{fig:yield}, we compare the estimated element masses for 0103$-$726 and 0045$-$734 with the theoretical predictions of chemical compositions produced with supernovae (SNe) from various progenitor masses \\citep{tsu1995}. Type-Ia origin would be rejected for both of the SNRs, because the data largely exceed the theoretical prediction for Ne and Mg, while the data do not show a large excess of Fe, as expected from theory. On the other hand, our data roughly agree with the prediction for a type-II SN case with a progenitor mass of $\\lesssim 20 M_{\\solar}$. Since these SNRs are located in H~{\\sc ii} regions, the type-II SNe origin is reasonable.   \\begin{figure}[hbtp] \\begin{center} \\hspace*{4mm} \\psfig{file=figure3a.ps,width=0.46\\textwidth,clip=}  \\hspace*{4mm} \\psfig{file=figure3b.ps,width=0.46\\textwidth,clip=} \\caption{Nucleosynthesis products of various SNe overlaid with the masses of elements in (a) 0103$-$726 and (b) 0045$-$734 derived from the spectral analyses. Solid lines with asterisks and diamonds represent products in type-II and type-Ia SNe, respectively \\citep{tsu1995}. The progenitor masses of the type-II SNe are also indicated. The dotted lines with circles, triangles, and squares represent models I, II, and III, respectively. Model II for 0103$-$726 and model III for 0045$-$734 are not plotted because of the bad fitting (see subsection \\ref{subsec:spec}). Plots for the elements fixed to be 0.2 solar in the fitting are not presented. } \\label{fig:yield} \\end{center} \\end{figure}   The spectrum of 0057$-$7226 does not exhibit overabundances, and is consistent with the SMC ISM. Therefore, the type of the SN cannot be constrained. However, its location in a giant H~{\\sc ii} region prefers a type-II origin.  Since the total mass of the ``metal-poor'' SNR 0057$-$7226 is comparable to those of ``metal-rich'' SNRs 0103$-$726 and  0045$-$734,  0057$-$7226 may simply have a smaller ejecta mass than the other two SNRs. One may argue, alternatively, if the progenitor is less massive than those of 0103$-$726 and 0045$-$734, for example   $\\sim 13 M_{\\solar}$, that the abundances of O, Ne, and Mg may be explained (see solid lines in figure \\ref{fig:yield}). However, this model predicts a larger mass of  Fe, which is not found in the observed spectrum.   \\subsection{Implication for the SNR Evolution}  We have shown that the X-ray emissions of the three SNRs are predominantly concentrated inside the radio shells (hereafter, ``centrally peaked'' morphology). The X-ray spectra predict that the ages are all rather old ($\\gtrsim 10^4$~yr), while two of them are found to exhibit high abundances possibly due to the ejecta gas.  To explain these features as well as the results from the LMC SNRs (\\cite{wil1999}; \\cite{hug1998}; \\cite{nis2001}), we  propose the following scenario. Soon after an SNR enters the Sedov phase (i.e., middle-aged stage), X-ray emission from the swept-up ISM becomes dominant because of the higher density at the shell, and hence the SNR exhibits a shell-like morphology and a low-abundance spectrum. As the age increases, the abundances gradually decrease and approach the mean value of the ISM, due to dilution with the swept-up ISM. At the same time, the morphology gradually changes from shell-like (``Shell'' and ``Diffuse Face 1'' in \\cite{wil1999}) to centrally-peaked type (``Diffuse Face 2'' and ``Centrally Brightened'' in \\cite{wil1999}), due to rapid cooling of the shell compared with the inner region. X-rays from the central region may thus be fossil radiation (\\cite{rho1998}, and references therein) and could be enhanced by the evaporation of cloudlets \\citep{whi1991}. The X-ray-emitting plasma shows a scatter of metal abundances, depending  on the  ejecta mass (large ejecta mass for  0103$-$726 and 0045$-$734 and small ejecta mass for  0057$-$7226). This scenario would be supported by numerical simulations by \\citet{she1999}.  The middle-aged LMC SNRs analyzed by \\citet{hug1998} and \\citet{nis2001} contain five shell-like and  seven centrally-peaked SNRs \\citep{wil1999}; among them, one shell-like and three centrally-peaked SNRs were found to be overabundant. The fact that the ratio of the overabundant SNRs is higher for the centrally-peaked (3/7) than for the shell-like (1/5) may support the above scenario, although the statistics are still limited.  In a transition phase from a shell-like to a centrally-peaked SNR, a faint X-ray shell, like that of 0103$-$726, may be observed. In this case, the X-ray spectrum of the faint shell should exhibit low abundances consistent with the SMC ISM, while the inner part should exhibit an ejecta-dominated spectrum. Spatially resolved spectroscopy with Chandra or XMM-Newton should clarify this prediction. \\par \\vspace{1pc}\\par The authors are grateful to Dr.\\ K.\\ Yoshita, whose strong criticism and useful comments made the initial draft much better. The authors also thank constructive comments from the referee, Dr.\\ R.\\ Williams. J.Y.\\ and K.I.\\ are supported by the JSPS Research Fellowship for Young Scientists. We retrieved ROSAT data from the HEASARC Online System which is provided by NASA/GSFC."
    },
    "0201/astro-ph0201396_arXiv.txt": {
        "abstract": "Here we present a pilot study into whether elliptical galaxy counts alone, can place a useful constraint on the $\\Omega_M - \\Omega_{\\Lambda}$-plane. The elliptical galaxy counts are drawn from three surveys: The Millennium Galaxy Catalogue ($16 > B_{KPNO} > 20$), the B-band Parallel Survey ($20 > B_{AB} >24$) and the Hubble Deep Fields ($23 > B_{AB} > 28$). The elliptical luminosity function used in the modeling was derived from a combination of the Millennium Galaxy Catalogue, the two-degree field galaxy redshift survey and the Sloan Digital Sky Survey ($M_{*}^{E/S0}=-19.90, \\phi_{*}^{E/S0}=0.0019$ Mpc$^{-3}$ and $\\alpha^{E/S0}=-0.75$ for $H_{o}=75$km/s/Mpc). We adopt a benchmark model and tweak the various input parameters by their uncertainties to determine the impact upon the counts. We find that {\\it if} the faint-end slope of the elliptical galaxy luminosity function is known to $\\Delta \\alpha < 0.1$, then over the magnitude range $16 < B < 23$ the counts depend most critically upon the cosmology, and can be used to place a weak constraint on the $\\Omega_{M}-\\Omega_{\\Lambda}$-plane. ",
        "introduction": "Galaxy number-counts as a cosmological probe have been fraught with difficulties since the conception of the idea by Edwin Hubble in the 1930s (Hubble 1936).  In a recent review Sandage (1997) provides an insightful historical overview of this topic.  The concept itself was severely challenged by the work of Tinsley (1977) who showed that the faint galaxy number-counts depend more critically upon evolution than upon the cosmological model (at this time only zero-$\\Lambda$ models were being considered). In the 1990s even this was superseded by the faint blue galaxy problem (see Ellis 1997 for a recent review). It is therefore fair to say that in the lead up to the turn of the Millennium the use of galaxy number-counts as a cosmological probe was discredited. Nevertheless attempts were made and retrospectively may have provided the first tentative evidence for a positive cosmological constant (Yoshi \\& Peterson 1995).  Three factors make the possibility of a revival of galaxy number-counts credible, these are: morphological segregation of the faint galaxy population and in particular the extraction of ellipticals; advances in our understanding of the evolution (or rather non-evolution) of ellipticals (see summary by Peebles in these proceedings); and the advent of $\\Lambda$ which dramatically broadens the impact of the cosmology upon the counts. ",
        "conclusions": ""
    },
    "0201/astro-ph0201439_arXiv.txt": {
        "abstract": "I discuss recent work on the nature of relativistic winds from Rotation Powered Pulsars, on the physics of how they transport energy from the central compact object to the surrrounding world, and how that energy gets converted into observable synchrotron emission. ",
        "introduction": "Rotation Powered Pulsars (RPPs) are the prime examples where the electromagnetic extraction of energy from a rotating compact object clearly provides the power for the nonthermal photon emission that distinctively characterizes the world of High Energy Astrophysics. While the fact that the energy extraction reflects electromagnetic torques exerted on a neutron star by the macroscopic electromagnetic fields with which each star is endowed - indeed, the magnitude of that torque allows one to infer the strength of each star's magnetic moment - has been known since the earliest days of pulsar research (Gold 1968), and indeed was predicted (Pacini 1967) even before pulsars were discovered, 1) the physics of the processes through which the extraction works, 2) the physics of how the rotational energy is transmitted to the surrounding world, and 3) the physics of how that energy transforms into the observed synchrotron radiation from the nebulae around pulsars (when these are observed) have all remained open questions through more than 30 years of pulsar research. Answers to all three questions are of significance not only to the understanding of RPPs themselves, but also to the physics of Active Galactic Nuclei, especially issues of jet formation and blazar emission, and to the workings of Gamma Ray Burst sources, especially to those issues involving the physics of relativistic shock waves and possibly to the nature of the basic engine underlying the GRB phenomenon.  The earliest theoretical answer to 1) and 2) was the vacuum wave theory of electromagnetic spindown, a theory which has never satisfactorily led to useful answers to 3).  Modern pulsar theory suggests the answer to 2) is that a RPP throws off its rotational energy in the form of a relatively dense magnetized, relativistic wind of plasma (Michel 1969), largely composed of free electron-positron pairs with an embedded, wound up magnetic field. Support for this idea has always come from X-ray astronomy. Observations of course provide the basic input to all attempts to model these phenomena. Recent advances  have been especially powerful drivers of progress on question 3) and to some aspects of question 2).  Study of the Crab Nebula, whose Chandra X-ray picture (Weisskopf {\\it et al.} 2000) appears in Figure 1, and more recently of other young pulsar wind nebulae (PWN), has been the source of most of the conceptual machinery in this field.  The most widely accepted hypothesis for understanding what we see is that the observed nebular emission is synchrotron radiation from particles heated into a nonthermal distribution of energies by the termination shock waves in the magnetized relativistic $e^\\pm $ pair outflow(s) from the underlying pulsar. My focus (obsession, if you like) is on the physics of these shock(s) and what that has to to tell us about the wind powering the shock. The observations (including optical observations first done by Lampland 1921) demonstrate the intrinsic time variability of this shock transition. Therefore, time dependent modeling of the dynamics of these relativistic flows is needed, with results which yield quite interesting constraints on the central engine, as well as telling us a lot of interesting things about the shocks themselves, conclusions of relevance to AGN and GRB as well as to pulsars and the PWN.  One important fact must be kept in mind in interpreting these rather different looking systems. The Crab is a relatively compact, high magnetic field nebula.  At X-ray emitting energies, the radiating particles lose their energy in approximately a flow time from the termination shock in the pairs (perhaps at or just inside the inner X-ray ring) to the outer edge of the visible X-ray torus, if the flow velocity is that of the Rees and Gunn (1974) model (recovered in the MHD model of Kennel and Coroniti 1984a), $ v = (c/3) (R_{shock}/r)^2 $. This strong cooling is behind the well know fact that the nebular image contracts with increasing photon energy from infrared through 10 keV X-rays - interestingly, what little is known of the nebular size at higher photon energies (Pelling {\\it et al.} 1987) suggests that nebular image contraction may come to an end in the 100 keV - 1 MeV range, which if true would suggest that the shock heating = ``acceleration'' is distributed over an observationally resolvable region.  Other nebulae are not always so radiatively efficient. G320.4 around PSR B1509-59 clearly has a much weaker magnetic field, doubtless reflecting different confinement properties of its local interstellar medium, with synchrotron losses affecting the spectra of outflowing particles only at much larger radii, if at all (Gaensler {\\it et al.} 2002), even at X-ray emitting energies. 3C 58 also seems to have rather low radiative efficiency.  The nebulae's appearance shows energy injected in equatorial outflow and polar jets. The ratio of polar to equatorial energy flux is not known, nor do we know whether the observed brightness distributions, including the apparent angular gaps between equatorial torii and polar jet, reflect gaps in the energy injection or changes in the energy conversion processes, such that mid latitude energy injection is darker than polar and equatorial phenomena. ",
        "conclusions": ""
    },
    "0201/astro-ph0201325_arXiv.txt": {
        "abstract": "As more cooling flow clusters of galaxies with central radio sources are observed with the {\\it Chandra} and {\\it XMM-Newton} X-ray Observatories, more examples of ``bubbles'' (low-emission regions in the X-ray coincident with radio emission) are being found. These bubbles are surrounded by bright shells of X-ray emission, and no evidence of current strong shocks has yet been found. Using an analytic approach and some simplifying assumptions, we derive expressions relating the size and location of a bubble, as well as the density contrast between the bubble and the ambient medium, with the shock history of the bubble.  These can be applied straightforwardly to new observations.  We find that existing observations are consistent with a mild shock occurring in the past, and with the bulk of the cool material in the X-ray shells being cooled at the cluster center and then pushed outward by the radio source.  Strong shocks are ruled out unless they occurred more than 1 Gyr ago.  We also discuss Rayleigh-Taylor instabilities as well as the case of a bubble expanding into an older bubble produced from a previous cycle of radio activity. ",
        "introduction": "\\label{sec:intro}  {\\it Chandra} high spatial resolution observations of clusters of galaxies reveal the presence of X-ray-deficient bubbles in the inner regions of many cooling flow clusters (Perseus [Abell~426], Fabian et al.\\ 2000, 2002; Hydra A [Abell~780], McNamara et al.\\ 2000; Abell~2052, Blanton et al.\\ 2001, hereafter BSMW; Abell~496, Dupke \\& White 2001; Abell~2199, Fabian 2002; MKW3s, Mazzotta et al.\\ 2001; Abell~2597, McNamara et al.\\ 2001; RBS797, Schindler et al.\\ 2001; Abell~85, Kempner, Sarazin, \\& Ricker 2002; Abell~133, Fujita et al.\\ 2002; Abell~4059, Heinz et al.\\ 2002). Earlier {\\it ROSAT} observations also showed a few similar results, albeit with less resolution (Perseus [Abell~426], B\\\"ohringer et al.\\ 1993; Abell~4059, Huang \\& Sarazin 1998). These bubbles are characterized by low X-ray emissivity implying low density. In most cases, the bubbles are sites of strong radio emission. There are clusters where the bubbles are less well-defined, although there may be hints to their existence (Abell~1795, Fabian et al.\\ 2001; 3C295, Allen et al.\\ 2001b; Virgo M87, Belsole et al.\\ 2001). The absence of evidence for shocks suggests that the bubbles are expanding and moving at subsonic or mildly transonic velocities (Fabian et al.\\ 2000; McNamara et al.\\ 2000; BSMW). The presence of bubbles which do not coincide with strong radio emission (known as `ghost bubbles' or `ghost cavities') located farther from the centers of the clusters in Perseus (Fabian et al.\\ 2000), MKW3s (Mazzotta et al.\\ 2001), and Abell~2597 (McNamara et al.\\ 2001), suggests that the bubbles rise buoyantly.  The discovery of bubbles and their detailed observational study has stimulated theoretical studies of X-ray bubble formation and evolution (Churazov et al.\\ 2001; Nulsen et al.\\ 2002, hereafter N2002; Fabian et al.\\ 2002). N2002 study the origin of the cool gas in the rims of enhanced X-ray emission in Hydra A, and particularly address the role of magnetic fields. Although the current paper addresses some similar questions to those addressed by N2002, our analysis and results have only a small overlap with those of N2002. Our discussion is not specific to an individual cluster. We aim to provide simple analytical expressions that can be used for different conditions and evolutionary stages in a variety of cooling flow clusters. We illustrate our results by applying them to several specific clusters. A number of recent papers provide detailed numerical simulations of bubbles (Rizza et al.\\ 2000; Churazov et al.\\ 2001; Br\\\"uggen et al.\\ 2002; Quilis, Bower, \\& Balogh 2001; Saxton, Sutherland, \\& Bicknell 2001; Reynolds, Heinz, \\& Begelman 2001; Br\\\"uggen \\& Kaiser 2001). However, it is difficult to generalize the results of these simulations and apply them to other observed cooling flow clusters.  In \\S~\\ref{sec:shock}, we briefly discuss the thermal evolution of shocked gas near the cluster center. Simple expressions for the properties of a bubble inflated by hot plasma (e.g., from a radio jet) are derived in \\S~\\ref{sec:bubbles}. In \\S~\\ref{sec:rt}, we show that a bubble is stable to Rayleigh-Taylor modes at early stages, becoming unstable only at late stages and in the outer (away from the cluster center) portion. Some properties of ``ghost'' bubbles (i.e., bubbles at relatively large radii with weak or no radio emission) are discussed in \\S~\\ref{sec:ghost}. In \\S~\\ref{sec:interaction}, we consider the interaction between bubbles, which we speculate may explain the X-ray and radio structure in the eastern radio `ear' in M87. We discuss and summarize our main results in \\S~\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  We derived simple expressions which can be used to analyze some properties of X-ray deficient bubbles in clusters of galaxies. We demonstrated their application for Abell~2052 and Perseus. We also addressed the X-ray and radio structure to the east of the center of M87.  Our main results can be summarized as follows:  (1) The cool material which forms the X-ray bright shells around the X-ray deficient radio bubbles in the cooling flow cluster Abell~2052, and very likely in other cooling flow clusters (e.g., Perseus: Fabian et al. 2002, Hydra A: N2002), was lifted from the very inner regions of the cooling flow, rather than being shocked and then cooled substantially due to shock compression. The material went through a shock, but a mild one, the main effect of which was to shape the material in a dense shell around the radio bubble. To achieve the present shell-to-ambient density ratio of $\\sim 2$, the dense gas had to have an initial density $\\sim 4$ times its present ambient density. This is quite plausible for dense gas lifted from $\\sim 5-10$ kpc. One problem with this model is that the observed mass in the shells in Abell~2052 is somewhat larger than the expected mass in these inner regions, determined by extrapolating the observed gas density profile into the center. We suggest that a steeper density gradient in the inner region prior to the bubble inflation, and/or inhomogeneities of the shell and ambient medium can account for this mass discrepancy in our simple model.   (2) The presence of cool material in the X-ray shells surrounding the bubbles has implications for a possible strong AGN shock propagating from the cluster center at the onset of the radio activity. Such a shock is not observed presently in any cluster, but may have a duty cycle on a long time scale. Soker et al.\\ (2001) proposed that such strong intermittent shocks, with shock velocities $\\gtrsim 7000$ km s$^{-1}$, may result in much lower mass cooling rates. If, for example, the presently cool shell was at an initial temperature of $\\sim 10^{7}$ K, and then was hit by a strong shock with a speed of $\\sim 7000$ km s$^{-1}$ (a Mach number ${\\cal{M}} \\sim 15$). From equations~(\\ref{eq:shock compression}), (\\ref{eq:shock pressure}), and (\\ref{eq:max_cool_f}), we find that the gas was shocked to a temperature of $\\sim 6 \\times 10^8$ K, and then cooled adiabatically (the radiative cooling will be negligible during a time of $< 10^{8} \\, {\\rm yr}$) to $\\sim 7 \\times 10^7$ K. This is still too hot to match the presently observed dense shell. We therefore rule out such a scenario. Cases with even higher Mach numbers were ruled out in $\\S 2$. We conclude that the presence of massive dense and cool shells preclude that the radio activity was initiated by a very strong shock (if the age of the present radio activity is $\\lesssim 10^8$ yr). However, this does not rule out a strong shock $\\gtrsim 10^9$ yr ago.  (3) From observations, it appears that the inflation of bubbles by AGN activity in cooling flow clusters is a general phenomenon. Our estimate via simple approach of the bubble size explains this as a result of the insensitivity of the bubble size (equations~\\ref{eq:bubble_radius1} and \\ref{eq:bubble_radius2}) to the injection energy rate and ambient density. Therefore, a bubble will be formed with a size of several kpc for most reasonable levels of AGN activity.  (4) During the early inflated phase, i.e., while energy is injected by the AGN, the interface between the hot bubble and the dense shell is Rayleigh-Taylor (RT) stable, explaining the smooth-surface bubbles close to the center of cooling flow clusters (e.g., Perseus). Later, the upper segment of the shell becomes unstable, and the shell may break-up there, e.g., the northern shell in Abell~2052. The interface between the dense shell and the outer lower density medium is RT-unstable, but the density ratio is close to unity, and considering some stress and friction, the instability evolves slowly. It is quite possible that the dense shell will start to break-up by these later RT-unstable modes. After being buoyant to about twice their diameter, the interface between the low density bubbles and the outer regions becomes RT-unstable, and large RT-fingers with sizes not much smaller than the bubble size may be observed penetrating the bubbles. The outer bubbles in Perseus seem to be disrupted by RT instabilities.  (5) We study the case of a bubble expanding into an older bubble. The portion of the still-active bubble which enters the low density cavity formed by the older bubble will be accelerated to higher velocities, up to $\\sim 2$ times the original expansion velocity of the bubble. These quickly moving blobs may travel out to radii of $20-30$ kpc in reasonable times. We speculate that such high density blobs may explain the high density X-ray emitting regions along the radio structures in M87. In M87, the cavity was formed by a radio jet which formed the eastern `ear', and the blobs created a stream of dense X-ray emitting gas."
    },
    "0201/astro-ph0201055_arXiv.txt": {
        "abstract": "We study the growth rate of stars via stellar collisions in dense star clusters, calibrating our analytic calculations with direct \\nbody\\, simulations of up to 65536 stars, performed on the GRAPE family of special-purpose computers. We find that star clusters with initial half-mass relaxation times $\\aplt 25$\\,Myr are dominated by stellar collisions, the first collisions occurring at or near the point of core collapse, which is driven by the segregation of the most massive stars to the cluster center, where they end up in hard binaries. The majority of collisions occur with the same star, resulting in the runaway growth of a supermassive object.  This object can grow up to $\\sim0.1$\\% of the mass of the entire star cluster and could manifest itself as an intermediate-mass black hole (IMBH). The phase of runaway growth lasts until mass loss by stellar evolution arrests core collapse.  Star clusters older than about 5\\,Myr and with present-day half-mass relaxation times $\\aplt 100$\\,Myr are expected to contain an IMBH. ",
        "introduction": "Using the {\\em Chandra} X-ray observatory, Kaaret et al.\\, (2000; 2001)\\nocite{2001MNRAS.321L..29K}\\nocite{2000HEAD...32.1511K} and Matsumoto et al.\\, (2000; 2001)\\nocite{2000HEAD...32.0108M}\\nocite{2001ApJ...547L..25M} recently discovered nine bright X-ray sources in the irregular galaxy M82. Their brightest source (No.\\ 7 in Table\\,1 of Matsumoto et al.\\, 2001) has a luminosity of $9 \\times 10^{40} {\\rm erg\\, s}^{-1}$ in the 0.2--10\\,KeV band, corresponding to the Eddington luminosity of a $\\sim 600$\\,\\msun\\, compact object.  The high luminosity and rather soft X-ray spectrum of the object indicates that it may be an intermediate-mass black hole (IMBH) with a mass of at least 600\\,\\msun\\, (Kaaret et al.\\, 2001; Matsumoto et al.\\, 2001).  An optical follow-up in the infrared (J, H, and K$^\\prime$-bands) with the CISCO instrument on the SUBARU telescope revealed a star cluster with an estimated mass of a few $10^6$\\,\\msun\\, at a position consistent with the X-ray location of the IMBH (Harashima et~al.\\, 2001).  This star cluster appears to be very young ($\\aplt 10$\\,Myr), as it is extremely blue and expanding shells of molecular gas have been discovered in its vicinity (Matsushita et al.\\, 2000),\\nocite{2000ApJ...545L.107M} typical for a star-forming region of a few million years.  Matsushita et al.\\, (2000) estimate that the environment has an age of only a few million years.  More unusually bright X-ray point sources have been discovered in the early spiral galaxies NGC 2403 (Kotoku et al.\\, 2000)\\nocite{2000PASJ...52.1081K} and NGC 4565 (Mizuno, et al. 1999)\\nocite{1999PASJ...51..663M}. Most remarkable, however, is the discovery of many bright X-ray sources in the ``Antennae'' system (NGC 4038/4039) by Fabbiano et al.\\,(2001)\\nocite{2001ApJ...554.1035F}, Zazas \\& Fabbiano (2002)\\nocite{2002astro.ph..3176Z} and Zazas et al.\\, (2002)\\nocite{2002astro.ph..3174Z}, also using {\\em Chandra}. These authors conclude that many of these sources may be $\\apgt 100$\\,\\msun\\, accreting black holes (although alternative explanations exist---see e.g.\\, Mizuno 1999; King et al.\\, 2001).\\nocite{2001ApJ...552L.109K}\\nocite{1999AN....320..356M} The Antennae contain many young star clusters with characteristics similar to those found in M82 (Mengel et al.\\, 2001).\\nocite{2001ApJ...550..280M} However, it is not yet clear how many of the X-ray sources in the Antennae are associated with these clusters (Zazas \\& Fabbiano\\, 2002).\\nocite{2002astro.ph..3176Z} There may also be an example of an IMBH in our own Galaxy, as recent reverberation mapping of the globular cluster M15 by Gebhardt et al.\\, (2000, 2001, and private communication)\\nocite{2000AJ....119.1268G}\\nocite{2001AAS...199.5610G} strongly suggests that the cluster may harbor a $\\sim 2500$\\,\\msun\\, black hole at its center.  Several possible mechanisms for forming IMBHs in star clusters have recently been suggested.  Miller \\& Hamilton (2001) have studied the possibility that an IMBH may form slowly (on a Hubble time scale) by occasionally encountering and devouring other cluster stars.  Mouri \\& Taniguchi (2002) have proposed a much more rapid black-hole merger mechanism, operating in very high-density ($10^6$ black holes pc$^{-3}$) environments on time scales as short as $\\sim10^7$ yr.  In this paper we consider the possibility of forming a massive object in a young star cluster due to repeated collisions during an early phase of core collapse.  Sanders (1970),\\nocite{1970ApJ...162..791S} Lee (1987),\\nocite{1987ApJ...319..801L} and Quinlan \\& Shapiro (1990)\\nocite{1990ApJ...356..483Q} have studied the possibility of collision runaways in spherical stellar systems of $\\apgt 10^7$ stars with high ($>100$\\kms) velocity dispersions.  All studies began with stars of equal masses and found that, for sufficiently high densities and velocity dispersions, runaway mergers could indeed occur.  Quinlan \\& Shapiro observed that the collision time scale for massive stars decreases faster with increasing mass than does the main-sequence lifetime, and concluded that clusters with initial relaxation times of 1--$5 \\times 10^8$ years could grow a massive $\\apgt 100$\\,{\\msun} object by multiple mergers.  Sanders' (1970) Monte-Carlo calculations neglected the effects of mass segregation and found collision runaways only after mergers had driven the cluster into a state of high central density.  However, in the self-consistent Fokker--Planck models of Lee and Quinlan \\& Shapiro, the runaway started well before core collapse occurred.  All authors concluded that runaways would not occur in clusters containing less than $\\sim10^6-10^7$ stars because three-body binary heating in small $N$-systems provided sufficient energy to reverse core collapse before the runaway process could begin.  In contrast to the studies just described, the models discussed in this paper begin with a broad range of stellar masses.  Vishniac (1978)\\nocite{1978ApJ...223..986V} demonstrated that a Salpeter (1955)\\nocite{1955ApJ...121..161S} initial mass function is Spitzer (1969) unstable.\\nocite{1969ApJ...158L.139S} As a result, young star clusters may experience core collapse on the time scale on which the most massive stars segregate to the cluster center.  This time scale may be much shorter than the main sequence lifetimes of the stars involved.  Vishniac suggested that such a prompt collapse might lead to the formation of a massive compact object.  We find that early core collapse in a relatively low-$N$ star cluster may result in a collision runaway, so long as the most massive stars remain on the main sequence while the collapse occurs.  The possibility of multiple collisions involving the same star in a dense star cluster was demonstrated convincingly by Portegies Zwart et al. (1999), using the special-purpose GRAPE-4 (Makino et al.~1997)\\nocite{1997ApJ...480..432M} to speed up their direct N-body calculations with up to 12288 stars.  They concluded that, even in small clusters, runaway collisions may lead to the growth of a single massive star.  The earlier arguments that three-body binary heating would drive the expansion of the cluster core appear to be unimportant in these simulations, as mergers between stars tend to destroy binaries before they can heat the cluster effectively.  Indeed, in contrast to the underlying assumptions of previous collision studies, dynamically formed binaries in dense clusters act as a {\\em catalyst} for stellar mergers, boosting the collision rate far beyond the simple two-body expressions used in earlier work.  The N-body simulations covered a rather limited part of the available parameter space, but the initial conditions were selected to mimic known dense star clusters in the Galaxy and the Large Magellanic Cloud.  If the bright X-ray source in M82 does indeed correspond to a compact object of $\\apgt 600$\\,\\msun, this IMBH could have been formed by a collision runaway resulting from collapse of the cluster core early in the cluster.  In fact, as we will see, it is quite natural to expect a $\\sim 10^3$\\,\\msun\\, black hole in a million-solar-mass star cluster. We begin by deriving (\\S2) some simple analytic expressions describing the dynamical behavior observed in cluster simulations.  In \\S3\\, we calibrate these relations using direct \\nbody\\, simulations.  We then (\\S4) extend these results to derive simple relations between black-hole formation and cluster parameters. ",
        "conclusions": "We study the runaway growth of a single star in a dense star cluster using a combination of complementary approaches.  Our semi-analytic analysis is supported by detailed \\nbody\\, calculations in which the effects of stellar evolution, stellar dynamics, binary evolution and the perturbing effect of a background Galactic potential are taken self-consistently into account.  Star clusters with initial half-mass relaxation times $\\trlx \\aplt 25$\\,Myr experience a phase of runaway growth.  In this phase a single seed star grows to a mass of about 0.1\\% of the total mass of the cluster.  The first collision occurs at the moment the cluster core collapses. This happens at about 0.2\\,\\trlx\\, but no later than about 5\\,Myr (the evolution time scale for a $\\apgt 50\\msun$ star).  The star which experiences the first collision becomes the target for further collisions, initiating runaway growth.  The growth phase is terminated by (1) the disruption of the cluster in the tidal field of the Galaxy (at $t\\aplt 5\\trlx$) or (2) the reversal of core collapse by mass loss from the evolving stellar population (after about 25\\,Myr).  A star cluster can survive for longer than $5\\,\\trlx$ if, for example, it did not initially fill its Jacobi surface (``Roche lobe'') in the Galactic tidal field.  (Examples are NGC\\,3603 and R\\,136, the dense star cluster in the 30 Doradus region in the Large Magellanic cloud.) Such clusters go though a phase of runaway stellar growth, but recover after stellar mass loss drives the re-expansion of the cluster core.  >From an observational point of view, a tidally limited cluster experiences three very distinct evolutionary phases: a pre-collapse phase until 0.2\\,\\trlx, a phase of deep core collapse (from 0.2\\,\\trlx\\, to about 25\\,Myr), followed by an expansion phase eventually leading to the disruption of the cluster.  During the expansion phase the cluster half-mass radius expands, causing the mean relaxation time to increase by a factor of 4 (see Portegies Zwart et al.\\, 2001).  A cluster in this final phase will be observable with a current relaxation time less than $\\sim4\\times25\\,{\\rm Myr} = 100$\\,Myr. The clearest indication of its previous phase of core collapse and runaway growth would be the presence of a central compact object with a mass $\\aplt 0.1$\\% of the initial cluster mass.  The cluster may also be relatively depleted in low-mass compact objects (stellar mass black holes and neutron stars), as these are consumed during the runaway growth phase.  Star clusters with an initial relaxation time $\\trlx \\apgt 25$\\,Myr do not experience a phase of runaway growth, as core collapse is prevented by mass loss from the most massive stars. These clusters may experience core collapse after $\\sim 100$\\,Myr, when stellar evolution slows (Takahashi \\& Portegies Zwart 1999).\\nocite{1998ApJ...503L..49T} This later core collapse, however, does not lead to a phase of runaway growth. In such old clusters multiple collisions are still likely to be common and may lead to blue stragglers with a mass more than twice the turn-off mass."
    },
    "0201/astro-ph0201263_arXiv.txt": {
        "abstract": "Using a simplified model of a black hole-accretion disk system which is dominated by Poynting flux, the evolution of the central black hole which is supposed to be powering GRB is discussed.  It is demonstrated explicitly that there is a lower limit on the angular momentum parameter for a given GRB energy.  It is found that the most energetic GRBs can only accommodate relatively rapid-rotating black holes at the center.   For a set of GRBs for which  the isotropic energies and $T_{90}$s are known,   the effect of the disk mass and the magnetic field on the horizon are discussed quantitatively. It is found that the magnetic field has little influence on the energy but affects the GRB duration as expected. The role of the disk mass is found to be significant in determining  both the energy and the duration.  \\vskip .3cm  \\noindent PACS numbers: 97.10.Gz, 97.60.Lf ",
        "introduction": "The discovery and the recent observations followed afterwards  indicate that the central region  powering gamma ray burst is rather compact in size, less than $10^8$ cm,  from the studies of energies and temporal structures\\cite{piran}.    Merging binary compact objects, hypernovae and rotating black holes have been considered to be among the viable candidates for a GRB central engine.  The common feature of these models is the formation or the existence of a black hole at the central region.  However, the gamma ray bursts or the afterglows followed do not  provide any direct observational informations on the central region and we do not know the exact nature of  the black hole at the center.  Hence although it is very important and interesting to investigate the central object itself,  we should rely on the indirect method of using a particular model, in which the physical properties  of black holes can be inferred from the observational data.   Among the models proposed so far, we choose a model in which the rotational energy of the black hole is responsible for powering GRB\\cite{lwb}. The mechanism of tapping rotational energy  from the black hole has been known as Blandford-Znajek mechanism\\cite{bz}\\cite{TPM}, in which the rotational energy is extracted out to the loading region via the magnetic flux which threads on the horizon.  It is easy to demonstrate that the strong magnetic field of $\\sim10^{15}$gauss is consistent with the essential features of GRBs\\cite{lwb}.  It is well known that the  black hole itself cannot keep the magnetic field on the horizon. It disappears very rapidly if not the environment keeps them from disappearing.  Recently, it is shown that the magnetic flux on the strongly magnetized object can be maintained during the collapsing process together with the electric charge onto the black hole\\cite{llp}. When the gravitationally unstable  object collapses into black hole the most natural environment is the accretion disk/torus which emerges together with the central black hole and it can provide  the magnetoshpere which keeps the magnetic flux from disappearing. Because of the accretion the energy and the angular momentum  are carried into the black hole while the Poynting flux carries away energy and angular momentum out of the black hole\\cite{lbw}, the evolution of a black hole depends not only on the Poynting flux but also on the accretion.  In this work we make an attempt to infer the evolution of the central black hole   during the gamma rays  are bursting using a simple model of a black hole-accretion disk system suggested in \\cite{lk}, in  which the accretion is dominated by the Poynting flux\\cite{blandford}\\cite{hkl}. In section II, the simple minded model is sketched with emphasis on the parameters which govern the evolution of the black hole. We take the initial values of the magnetic field on the horizon and the disk mass as parameters while the black hole mass is initially taken to be a typical mass, $7M_{\\odot}$, suggested from the  black hole binary  systems. The evolution of the angular momentum parameter $\\tilde{a}$ is discussed in detail in section III. With finite mass of accretion disk, $M_D \\leq 7M_{\\odot}$, the initially rapid-rotating black hole is expected to remain rotating even at the end of GRB. But the initially slow-rotating black hole is found to eventually stop rotating with GRB.  In section IV, the total energy out of the rotating black hole is calculated and compared to  GRBs for which the isotropic energies are known\\cite{grb}. The lower limits of the angular momentum parameter of the black holes for the corresponding GRBs are calculated. For the most energetic GRBs, $E_{iso} > .1 M_{\\odot}$, the lower limit is found to be rather high, $\\tilde{a}(0) >0.3$.  In section V, the evolution of the black hole is discussed for GRBs with known $T_{90}$.  Identifying $E_{iso}$ as 90\\% of energy extraction and  $T_{90}$ as the time taken to extract out $E_{iso}$, the corresponding sets of $B_H(0)$ and $M_D(0)$ are determined to discuss the evolution of the corresponding black holes.  The discussions are given in section VI. ",
        "conclusions": "Using a schematic model for a central engine of GRB, which consists of a black hole and an accretion disk\\cite{lk} at the center, the evolution of the central black hole is discussed.  The accretion is assumed to be dominated by the Poynting flux out of the disk and the GRB is supposed to be powered by the Poynting flux which extracts out a part of the rotational energy of a black hole. It is found that the evolution of the rotation parameterized by $\\tilde{a}(t)$ shows different patterns depending on  the initial value, $\\tilde{a}(0)$.   For a given disk mass, the black hole with  $\\tilde{a}(0) < .57$  much more rapidly approaches to non-rotating black hole than with  $\\tilde{a}(0) > .57$.   It is also shown that there is a maximum energy for GRB for a given $\\tilde{a}(0)$. Hence one can infer the lower limit of the angular momentum parameter for the central black hole.  The most energetic GRBs, are found to be able to accommodate only rapidly rotating black holes , $\\tilde{a}(0)> 0.4$.  The effect of the magnetic field on the total energy for GRB  is found to be not significant compared to the disk mass.  However since the stronger magnetic field extract  energy more rapidly than the weaker magnetic field, the detailed variation of the GRB duration is found to be due to  the magnetic field structure of the system. The role of the disk mass in this model is significant both in determining the energy and the duration of GRB.  Within the range of the parameters  used to fit a set of GRBs for which isotropic energy and $T_{90}$ are well determined, the final black holes are found to become  more massive  than the initial values but with  smaller angular momentum parameters.   This observation is consistent with the  general feature expected for a system of  a black hole - accretion disk with finite size and life time. As it is mentioned the analysis is limited to GRBs with $T_{90}$ determined in this work. But provided with a systematic way of determining GRB duration time from the observational data, statistically the more meaningful conclusion can be made on the evolution of the black hole at the center of GRB and it remains as a future work. The characteristic feature of this model is the sign factor $A$ determined by the magnetically dominated accretion disk. However  it depends on the several simplifications which are subject to be verified. For example, we use the relation of the field components suggested by  Blandford which requires a justification if it can be used  in the relativistic formulation especially  for a rapidly rotating black hole.  Also the identification of the angular velocity of the field line as the Keplerian angular velocity in this model also needs a valid justification\\cite{hkl}.  \\vskip 0.5cm This work was supported by Hanyang University, Korea made in the program year of 2001.  \\newpage"
    },
    "0201/astro-ph0201096_arXiv.txt": {
        "abstract": "{ The origin of subluminous B stars is still an unsolved problem in stellar evolution. Single star as well as close binary evolution scenarios have been invoked but until now have met with little success. We have carried out a small survey of spectroscopic binary candidates (19 systems consisting of an sdB star and late type companion) with the Planetary Camera of the WFPC2 onboard Hubble Space Telescope to test these scenarios. Monte Carlo simulations indicate that by imaging the programme stars in the R-band about one third of the sample (6--7 stars) should be resolved at a limiting angular resolution of 0\\bsec1 if they have linear separations like main sequence stars (``single star evolution''). None should be resolvable if all systems were produced by close binary evolution. In addition we expect three triple systems to be present in our sample. Most of these, if not all, should be resolvable. Components were resolved in 6 systems with separations between 0\\bsec2 and 4\\bsec5. However, only in the two systems TON~139 and PG~1718$+$519 (separations 0\\bsec32 and 0\\bsec24, respectively) do the magnitudes of the resolved components match the expectations from the deconvolution of the spectral energy distribution. These two stars could be physical binaries whereas in the other cases the nearby star may be a chance projection or a third component. Radial velocity measurements indicate that the resolved system TON~139 is a triple system, with the sdB having a close companion that does not contribute detectably to the integrated light of the system. Radial velocity information for the second resolved system, PG~1718$+$519, is insufficient. Assuming that it is not a triple system, it would be the only resolved system in our sample. Accordingly the success rate would be only 5\\% which is clearly {\\it below} the prediction for single star evolution. We conclude that the distribution of separations of sdB binaries deviates strongly from that of normal stars. Our results add further evidence that close binary evolution is fundamental for the evolution of sdB stars. ",
        "introduction": "}  Subluminous B (sdB) stars dominate the populations of faint blue stars of our own Galaxy and are found in both the disk (field sdBs) and globular clusters (Moehler et al. \\cite{mohe97}). Observations of elliptical galaxies with the Ultraviolet Imaging Telescope (Brown et al.\\ \\cite{brfe97}) and the Hubble Space Telescope (Brown et al.\\ \\cite{brbo00}) have shown  that these stars are sufficiently common to be the dominant source for the ``UV upturn phenomenon'' observed in elliptical galaxies and galaxy bulges (see also Greggio \\& Renzini \\cite{grre90}, \\cite{grre99}). Their space distribution and kinematical properties indicate that the field stars belong to the intermediate to old disk population (de Boer et al. \\cite{deag97}; Altmann \\& de Boer \\cite{alde00}).  However, important questions remain concerning their formation process and the appropriate evolutionary timescales. This is a major drawback for the calibration of the observed ultraviolet upturn in elliptical galaxies as an age indicator.  It is now generally accepted that the sdB stars can be identified with models for Extreme Horizontal Branch (EHB) stars burning He in their core, but with a very tiny ($<$2\\% by mass) inert hydrogen envelope (Heber \\cite{hebe86}; Saffer et al.\\ \\cite{sabe94}). An EHB star bears great resemblance to a helium main-sequence star of half a solar mass and its further evolution should proceed similarly (i.e. directly to the white dwarf graveyard) as confirmed by evolutionary calculations (Dorman et al. \\cite{doro93}).  How stars evolve to the EHB configuration is controversial. The problem is how the mass loss mechanism in the progenitor manages to remove all but a tiny fraction of the hydrogen envelope at {\\em precisely} the same time as the He core has attained the minimum mass ($\\approx0.5$M$_\\odot$) required for the He flash.  Both non-interacting (scenario i), and interacting (scenarios ii and iii) evolutionary scenarios have been proposed to explain the origin of the sdB stars (see Bailyn et al. \\cite{basa92}).  \\noindent (i) Enhanced mass loss on the red giant branch (RGB) before or during the core helium flash may remove almost the entire hydrogen-rich envelope. This is usually modelled by increasing the $\\eta$ factor in the Reimers (\\cite{reim75}) formula to estimate mass loss rates for RGB stars. It has been conjectured that the mass loss rates increase with increasing metallicity, implying that metal rich populations should produce more sdB stars than metal poor ones. Birthrate estimates for sdB stars indicate that only 2\\% (Heber \\cite{hebe86}) or even less (0.25\\% to 1\\%, Saffer \\& Liebert \\cite{sali95}) of the RGB stars need to experience such enhanced mass loss. Evidence that this is possible comes from the existence of RR Lyrae stars of population I which must also have lost half of their mass during evolution. In both cases the physical reason for such strong mass loss is not yet understood.  \\noindent (ii) Mengel et al. (\\cite{meno76}) suggest that sdBs could be formed from binaries in which mass transfer starts on the red giant branch and results in a reduction of the hydrogen envelope prior to the helium core flash. Hence all sdBs star are predicted to be found in close binary systems.  \\noindent (iii) An alternative scenario was proposed by Iben (\\cite{iben90}), who pointed out that sdBs can be formed from mergers of helium white dwarf binary systems. Iben \\& Tutukov (\\cite{ibtu92}) estimate that 80\\% of the sdBs could have been formed by mergers. Hence the frequency of sdBs still being in binaries should be at most 20\\%.  Several dozens of objects with composite spectra consisting of an sdB and a dwarf G-K star have been discovered (e.g. Ferguson et al. \\cite{fegr84}; Theissen et al. \\cite{thmo93}, \\cite{thmo95}; Allard et al. \\cite{alwe94}) which implies that the binary frequency of sdBs is 50\\% or more (Allard et al. \\cite{alwe94}). The observed large binary frequency rules out the merger scenario (iii) and we are left with scenarios (i) and (ii), i.e. either the sdB binaries are mostly wide systems that did not interact so that the sdB precursors have evolved independently from the companion (i), or they are close systems formed by interaction of the sdB precursor with the companion star (mass exchange, ii).  The high spatial resolution of the {\\it Planetary Camera} (PC) on board the {\\it Hubble Space Telescope} (HST) allows to perform a crucial test. As we will show in this paper, it should be possible to resolve a significant fraction of the known composite spectrum systems containing an sdB star if scenario (i) is correct, i.e. if the systems have a distribution of separations like normal main sequence binaries (Duquennoy \\& Mayor \\cite{duma91}). The interacting scenario (ii), however, predicts that all sdB stars reside in short period (P$\\le$100d) binaries and consequently none of the systems should be resolvable even with the PC. In order to measure their distribution of separations we have imaged 23 sdB binary candidates with the PC by taking advantage of the snap shot mode of HST observations. ",
        "conclusions": "In total we have resolved six systems out of a sample of 23 stars. Of those 23 stars, however, four do not really belong to the intended sample of sdB stars showing evidence for a cool companion: PG~1558$-$007, KPD~2215$+$5037, and PG~2259$+$134 show no photometric or spectroscopic evidence for a companion. The observed infrared excess can be explained by interstellar reddening rather than by a cool companion.  PG~1558$-$007 does have a resolved near by star (linear separation 1500\\,AU), which, however, is too faint to contribute detectably to the combined light in the R band. PG~0105$+$276 is a helium-rich sdO star (with two possible distant companions at 3700\\,AU and 4900\\,AU).  Of the remaining four resolved systems the nearby stars are in two case (TON~1281, HE~0430$-$2457) too faint to reproduce the photometric and/or spectroscopic observations of the stars.  Only in the two systems TON~139 and PG~1718$+$519 (separations 0\\bsec32 and 0\\bsec24, respectively) do the magnitudes of the resolved components match the expectations. These two stars could be physical binaries whereas in the other cases the nearby star may be a third component or a chance projection. Radial velocity measurements indicate, however, that the resolved system TON~139 is also triple.   Hence, the observed sdB binary sample was reduced to 19 objects with two bona-fide resolved systems, which have apparent separations of 0\\bsec24 and 0\\bsec32. From the numerical simulations we would expect to resolve six to seven systems if sdB stars have the same binary characteristics as normal stars, out of which one system is expected to have $a \\sin i>$1\\arcsec\\ and two should have separations between 0\\bsec1 and 0\\bsec2. The discrepancy becomes even more pronounced if one recalls that our photometric fit procedure tends to underestimate the brightness of the companion (and thus to overestimate the limiting angular separation that can still be resolved). In addition we expect three triple systems to be present in our sample. Most of these, if not all, should be resolvable. Such systems could explain some of the more distant companions as well as the radial velocity measurements of TON~139.  This success rate (1 resolved binary out of 19 candidates) is clearly {\\it below} the prediction of numerical simulations assuming single star evolution (about 30\\%), using the distribution of binary separations given by Duquennoy \\& Mayor (\\cite{duma91}). This indicates that the distribution of separations of sdB binaries strongly deviates from that of normal stars.  If, on the other hand, all sdB stars were produced by close binary evolution, none of the binary systems should have been resolved (even at the high spatial resolution of the WFPC2 camera). Our low success rate is thus closer to that predicted by the close binary evolutionary scenario. Recent radial velocity surveys (Saffer et al. \\cite{sagr01}; Maxted et al. \\cite{mahe01}) revealed that a large fraction of single-lined sdB stars are indeed close binaries with periods below 10 days. Our results could be explained if most of the programme stars were close binaries. Therefore, our study provides further evidence that close binary evolution indeed is fundamental to the evolution of sdB stars. A survey for radial velocity variations in all of our programme stars will be tale telling."
    },
    "0201/astro-ph0201119_arXiv.txt": {
        "abstract": "I 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.  My group uses semi-analytic models (SAMs) of galaxy formation, set within the CDM 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.  We have improved our 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 also discuss recent attempts to determine the emitted spectrum of high energy gamma rays from blazars such as Mrk 501 using the synchrotron self-Compton model and the observed X-rays, and note that our favorite SAM EBL plus the observed spectrum of Mrk 501 do {\\it not} imply unphysical upturns in the high energy emitted spectrum --- thus undermining recent claims of a crisis with drastic possible consequences such as breaking of Lorentz invariance.  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.\\footnote{This paper is an updated version of \\protect\\cite{PGam2000}.} ",
        "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 approaches, by integrating the stellar SEDs and dust modelling of \\cite{dev1,dev2} into the galaxy formation SAM 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 (e.g.~\\cite{papovich}) and damped Lyman-$\\alpha$ systems (SPF and \\cite{maller}) 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.  Section 4 presents the implications for $\\gamma$-ray attenuation, and \\S5 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": ""
    },
    "0201/astro-ph0201433_arXiv.txt": {
        "abstract": "We investigate antikaon condensation in compact star matter using a relativistic mean field model. Antikaon condensates make the equation of state softer resulting in a smaller maximum mass star compared to the case without condensate. It is found that the equation of state including antikaon condensates gives rise to a stable sequence of compact stars called the third family beyond the neutron star branch. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0201/hep-ph0201089_arXiv.txt": {
        "abstract": "Assuming that the observed deficit of solar neutrinos is due to the interaction of their transition magnetic moment with the solar magnetic field we derive the predictions for the forthcoming Borexino experiment. Three different model magnetic field profiles which give very good global fits of the currently available solar neutrino data are used. The expected signal at Borexino is significantly lower than those predicted by the LMA, LOW and VO neutrino oscillation solutions of the solar neutrino problem. It is similar to that of the SMA oscillation solution which, however, is strongly disfavoured by the Super-Kamiokande data on day and night spectra and zenith angle distribution of the events. Thus, the neutrino magnetic moment solution of the solar neutrino problem can be unambiguously distinguished from the currently favoured oscillation solutions at Borexino. ",
        "introduction": "If lepton flavour is not conserved, neutrinos must have flavour-off-diagonal (transition) magnetic moments, which applies to both Dirac and Majorana neutrinos. Under a transverse magnetic field, such magnetic moments will cause a simultaneous rotation of neutrino spin and flavour, spin-flavour precession \\cite{SV,VVO}. This precession can be resonantly enhanced in matter \\cite{LM,Akh1,rev}, very much similarly to the resonance amplification of neutrino oscillations, the MSW effect \\cite{MSW}.  The resonance spin-flavour precession (RSFP) of solar neutrinos due to the interaction of their transition magnetic moments with the solar magnetic field can account for the observed deficit of solar neutrinos. The conversion mechanism is neutrino energy dependent, which is a necessary feature to fit the data.  RSFP requires relatively large values of the neutrino transition magnetic moment, $\\mu_\\nu \\sim 10^{-11}\\mu_B$ for peak values of the solar magnetic field $B_0\\sim 100$ kG. Although such values of $\\mu_\\nu$ are not experimentally excluded, they are hard to achieve in the simplest extensions of the standard electroweak model. Still, the RSFP mechanism yields an excellent fit of all currently available solar neutrino data (see, e.g., \\mbox{[7 -- 13]} for recent analyses), typically even somewhat better than does the large mixing angle (LMA) oscillation solution, which is the best one among the oscillation solutions. In any case, in pursuit of the solution of the solar neutrino problem it is very important to test all non-standard hypotheses, and neutrino magnetic moment seems to be the most plausible alternative to neutrino oscillations.  As non-vanishing neutrino transition magnetic moments imply lepton flavour violation, they must be accompanied by the usual lepton flavour mixing. Thus RSFP should in general coexist with neutrino oscillations. It is quite possible, however, that the flavour mixing in the solar neutrino sector is too small to be of any relevance to the solar neutrino problem. This is our assumption in the present paper, i.e. we neglect neutrino oscillations and consider pure RSFP transitions. Small flavor mixing in the solar sector does not contradict the large mixing in the atmospheric neutrino sector -- the corresponding mixing  angles are independent parameters. In this connection, one can recall that the lepton mixing angle $\\theta_{13}$ probed in short-baseline reactor neutrino experiments is known to be small or vanishing \\cite{chooz} even though the ``atmospheric'' mixing angle $\\theta_{23}$ is large \\cite{atm}.  Unfortunately, the RSFP solution of the solar neutrino problem is difficult to establish experimentally. Except for predicting reduced detection rates of solar neutrinos (which the oscillation solutions also predict), it has mostly negative signatures: No time variations beyond the usual $1/R^2$ variation due to the eccentricity of the Earth's orbit (assuming that the strongest component of the solar magnetic field does not vary with time)\\footnote{There is a caveat here which we shall discuss in Section 3 -- strictly speaking, this is only true when the solar magnetic field is spherically symmetric.}; no day-night effect; no significant distortions of the solar neutrino spectrum in Super-Kamiokande and SNO experiments. One might therefore think that the RSFP solution of the solar neutrino problem can only be established if all the oscillation solutions are experimentally ruled out. Such a ``negative'' confirmation would hardly satisfy anyone.  In the present paper we show that in fact this is not the case: the RSFP predictions for the Borexino experiment are very different from those of neutrino oscillations, and different solutions of the solar neutrino problem can therefore be unambiguously distinguished experimentally. ",
        "conclusions": "The main goal of this work was to investigate whether RSFP can be distinguished from the oscillation solutions of the solar neutrino problem at Borexino. There are four main types of oscillation solutions of the solar neutrino problem, depending on the allowed values of the leptonic mixing angle $\\theta$ and neutrino mass squared difference $\\Delta m^2$: Large mixing angle (LMA), small mixing angle (SMA) and low--$\\Delta m^2$ (LOW) MSW solutions, and also vacuum oscillation (VO) solution (for recent discussions see, e.g., \\cite{Lisi2001,KS2001,BGGPG,GG,aliani,Gago}). The LMA, LOW and VO solutions all predict the average suppression of the event rate at Borexino by 35 -- 40\\%, whereas in the case of the SMA solution a suppression by about a factor of five is expected.  The main feature of the RSFP mechanism which can be exploited in order to distinguish it experimentally from neutrino oscillations is the peculiar shape of the energy dependence of the survival probability of solar neutrinos: At high energies it resembles the $\\nu_e$ survival probability of the LMA oscillation solution, whereas at low energies it is similar to that of the SMA solution. A mismatch in the results of the experiments sensitive to the high-energy and low-energy parts of the solar neutrino spectrum would therefore be an indication for RSFP.  As can be seen from Table III, the RSFP mechanism predicts the suppression of the event rate at Borexino by about a factor of three. The maximum allowed at 99\\% CL reduced rate is 0.62; this only marginally overlaps with the minimum allowed at 3$\\sigma$ reduced rate in the case of the LMA solution (0.58, see Table 7 of ref. \\cite{BGGPG}). Thus, the predictions of the RSFP and LMA solutions are more than 5$\\sigma$ away from each other and the probability of mistaking one for another is very low.  The minimum allowed at 3$\\sigma$ values of the reduced rate at Borexino in the case of LOW and VO solutions, 0.54 and 0.53 respectively \\cite{BGGPG}, are slightly lower than that for the LMA solution, so that there is a larger overlap with the 99\\% CL prediction of the RSFP. However, in these cases, too, one can easily discriminate between RSFP and the oscillation solutions. Indeed, in the case of the LOW  solution one expects a sizeable (up to 40\\%) day-night event rate difference at Borexino, while VO should lead to large seasonal variations beyond the usual $1/R^2$ dependence. No such effects are predicted by RSFP.  Our predictions for the reduced event rate at Borexino in the case of RSFP are slightly higher than those of the SMA oscillation solution, although there is a significant overlap between the predicted rates in these two cases. It should be noted, however, that the SMA solution is strongly disfavoured by the data on day and night spectra and zenith angle distributions of recoil electrons at Super-Kamiokande \\cite{SK}. We therefore conclude that Borexino will allow a clear discrimination between RSFP and currently favoured oscillation solutions of the solar neutrino problem. It should be noted that new dedicated low-energy solar neutrino experiments, which are widely discussed now \\cite{lownu}, should have a similar or even stronger discriminative power \\footnote{We thank M. Nakahata for pointing this out to us.}.  The RSFP mechanism may also lead to some specific effects, absent in the case of neutrino oscillation solutions. If the solar magnetic field is not axially symmetric, the rotation of the Sun can lead to a time variation of the signal with the period equal to the solar rotation period (about 28 Earth's days). Seasonal variations of the signal can also occur due to the inclination (by about $7^\\circ$) of the solar equatorial plane to the Earth's orbit, provided that the solar magnetic field depends on the polar angle $\\Theta$. This effect depends on the three-dimensional structure of the solar magnetic field. For the model profile of ref. \\cite{Mir1}, the transverse component $B_\\perp \\propto \\sin\\Theta$; since for solar neutrinos reaching the Earth $\\Theta=90^\\circ \\pm 7^\\circ$, one finds seasonal variations of less than $\\pm 1.5\\%$ for charged-current signals. In the case of neutrino detection through $\\nu e$ scattering (Super-Kamiokande, SNO and Borexino), these variations are further diluted by the neutral-current contribution to the event rates. Thus, the seasonal variations of this kind are probably too small to be observable.  Another possible signature of RSFP is an observable flux of $\\bar{\\nu}_e$ from the Sun if neutrinos, in addition to transition magnetic moments, have a sizeable flavour mixing ($\\theta \\aprge 0.1$) \\cite{antinu}. The flux of solar $\\bar{\\nu}_e$'s at the level of 1\\% of the $\\nu_e$ flux can, in principle, be detected at Borexino and SNO. However, these signatures depend on additional assumptions about $\\theta$ and the structure of the solar magnetic field, whereas our predictions for the Borexino detection rate are essentially model independent. The only possible model dependence is contained in the choice of the solar magnetic field profile, and this freedom is severely constrained by the requirement of fitting the available solar neutrino  data. As a result, the predictions for the Borexino event rate, though somewhat different for different profiles (see Table III), all fall below those for the LMA, LOW and VO solutions.  In conclusion, we have shown that the Borexino experiment will be able to unambiguously distinguish RSFP from the currently favoured oscillation solutions of the solar neutrino problem.  \\vspace{0.1cm} \\noindent  {\\em Acknowledgements.} We are grateful to M. Nakahata for useful correspondence. E.A. was supported by the Calouste Gulbenkian Foundation as a Gulbenkian Visiting Professor at Instituto Superior T\\'ecnico."
    },
    "0201/astro-ph0201355_arXiv.txt": {
        "abstract": "We present  ultra-high resolution  spectra for a set of nearby F-G-K stars on, or close to, the main sequence. The wavelength shifts of stellar lines relative to their laboratory wavelengths are measured for more than a thousand Fe {\\sc i} lines per star, finding a clear correlation with line depth. The observed patterns are  interpreted as convective blue-shifts that become more prominent for weaker lines, which are formed in deeper atmospheric layers. A morphological sequence with spectral type or effective temperature is apparent. Two K giant stars have also been studied. The velocity span between weak and strong lines  for these stars is larger  than for the dwarfs and subgiants of similar spectral types. Our results show that  convective wavelength shifts may seriously compromise the accuracy of absolute spectroscopic radial velocities, but  that an empirical correction may be applied to measured velocities. ",
        "introduction": "Convection in late-type stars penetrates into the photosphere, producing inhomogeneities. Temperature variations of up to a few hundred Kelvin are apparent in optical images of the Sun, and a velocity field of  several kilometers per second is observed in spatially resolved solar spectra. The warmer upflows appear as bright granules surrounded by narrower, cooler downflows. While present technology cannot  resolve stellar granulation, there is evidence  of its presence in all late-type stars. The temperature and velocity inhomogeneities produce absorption line profiles which are shifted and asymmetric, effects that are readily noticed in  ultra-high dispersion stellar spectra.  An extensive literature exists on solar line asymmetries and shifts (see, e.g., Neckel \\& Labs 1990; Asplund et al. 2000; Pierce \\& Lopresto 2000). Modern stellar studies were pioneered by Dravins (1974, 1985) and Gray (1980, 1981). Dravins (1999) provides a short review of recent work. Line asymmetries and line shifts are consequences of the same phenomenon and, therefore, give largely equivalent information. When the number of lines measured is limited, but the available features are mostly unblended, line asymmetries carry  more information. On the contrary, if the spectrum is heavily crowded, line shifts, which can be measured for many lines, are likely to provide more information. A limited wavelength coverage, largely dictated by the detector size and the scarcity of high-accuracy laboratory wavelengths, steered stellar studies toward the analysis of line asymmetries.  Line bisectors -- a measure of the line asymmetry -- vary smoothly across the HR diagram (Gray \\& Toner 1986; Dravins 1987; Gray \\& Nagel 1989; Dravins \\& Nordlund 1990). Other parameters, such as rotation, chemical composition, magnetic fields, or binarity, are likely to play a role, but the limited data available have prevented an in-depth study. A recent comprehensive work on the Fe {\\sc i} spectrum  (Nave et al. 1994) has provided accurate laboratory wavelengths for 9501 lines. These data and improvements in astronomical instrumentation make it feasible and practical to turn to line shifts as a complement to line bisectors.  Gravitational and convective wavelength shifts systematically affect absolute determinations of radial velocities. Other effects can also introduce systematic radial velocity errors, but they are expected to be  less important for most stars (Lindegren, Dravins, \\& Madsen 1999). Gravitational redshifts are proportional to the mass-to-radius ratio. For spectral lines formed in the photospheres of dwarf stars, the gravitational shift is the same for all lines, in the range between 0.4 and 2 km s$^{-1}$. When accurate parallaxes and photometry are available, comparison with evolutionary models allows us to determine dwarf masses and radii within 8 \\% and 6 \\% , respectively (Allende Prieto \\& Lambert 1999). It is therefore possible to estimate the gravitational shift within 10 \\%, yielding worst-case uncertainties of $\\sim 0.2$ km s$^{-1}$. Convective line shifts may vary with spectral type. The velocity difference between weak and strong lines is about $\\sim 0.6$ km s$^{-1}$ for the Sun (Allende Prieto \\& Garc\\'{\\i}a L\\'opez 1998) and probably more than 1 km s$^{-1}$ for the F5 subgiant  $\\alpha$ CMi (Procyon; Allende Prieto et al. 2002), but  no systematic study across the HR diagram has been published. Obviously, accurate spectroscopic studies of absolute radial velocities cannot afford to  neglect these shifts.  In this paper, we analyze optical spectra of several late-type stars obtained with the High Resolution Spectrograph (HRS) coupled to the Hobby-Eberly Telescope (HET). The high quality and large spectral coverage of the spectra allow us to measure line shifts for a large number of Fe {\\sc i} lines. Section \\ref{obs} describes the observations and the data. In Section \\ref{doit} we  describe the analysis, and in \\S \\ref{end} we discuss the results. ",
        "conclusions": "\\label{end}  The pattern of the  velocity shifts is similar for all the stars, and in qualitative agreement with solar results. The net convective blue shifts of the lines strengthen toward deeper photospheric layers, and therefore affect the weak lines more than the strong ones. Despite our limited sample, there is an indication of a smooth dependence of the average velocity shifts with spectral type. This effect  is clearer in our analysis of line shifts than in previous studies using bisectors. As argued in the introduction, this is likely the result of line shifts measurements being less affected by blending than line bisectors in late-G and K type stars.  Our preliminary results indicate that the trend of line shifts as a function of line strength can be determined as a function of spectral type and gravity. At least for late-type dwarfs, assuming the strongest lines in the spectrum are free from convective shifts (as is the case for the solar photosphere), it is possible to correct for convective wavelength shifts to within $\\simeq 0.2$ km s$^{-1}$.   The measured range of wavelength shifts for neutral iron lines is larger in $\\alpha$ CMi than in cooler stars of classes IV and V. The range of wavelength shifts is also larger,  reaching up to $\\sim$ 1 km s$^{-1}$,  for the K giants than for the G-K dwarfs and subgiants in our sample, which is not surprising. Although their lower effective temperatures provide less flux to transport, their lower atmospheric pressure implies larger convective velocities. Fig. 1 suggests that, once we account for the observational errors, there is little room for intrinsic line-to-line scatter in the  F-G-K stars sampled here. Convection is supposed to cease in main sequence stars with spectral types earlier than about F2, but other velocity fields must be present in those atmospheres, judging from the asymmetry of the spectral lines (e.g. Gray \\& Nagel 1989). Determining accurate absolute radial velocities demands an understanding of the wavelength shifts in the spectra of these stars too.  Solar observations  have shown differences in the center-to-limb variation of the granulation along the central meridian and the equator (e.g. Beckers \\& Taylor 1980; Rodr\\'{\\i}guez Hidalgo, Collados, \\& V\\'azquez 1992). Observations of a large number of stars may then show peculiarities at a given spectral type depending on the orientation of the rotational axis.  Stars with enhanced  magnetic fields are also expected to be peculiar in terms of observed wavelength shifts, as the enhanced magnetic fields may hinder the convective motions. Furthermore, line shifts may vary with time for stars exhibiting an analog of the solar 11 year cycle.  Observations of line shifts for a significant number of stars with high quality are required for a deeper understanding of granulation in stellar atmospheres, its relationship with other stellar phenomena,  and the role of surface convection in the structure and evolution of stars. Our first results show that measurements of convective line shifts are a must in order to derive accurate absolute radial velocities.  We thank the HET staff for their outstanding job making possible science observations with HRS since day one. The Hobby-Eberly Telescope is operated by McDonald Observatory on behalf of The University of Texas at Austin, the Pennsylvania State University, Stanford University, Ludwig-Maximilians-Universit\\\"at M\\\"unchen, and Georg-August-Universit\\\"at G\\\"ottingen. NSO/Kitt Peak FTS data used here were produced by NSF/NOAO. This research was supported in part by the NSF (grant AST-0086321)."
    },
    "0201/astro-ph0201163_arXiv.txt": {
        "abstract": "{ Using the Plateau de Bure Interferometer and IRAM 30m Telescope, we observed $\\lambda$2-3mm absorption lines of CS, SO, \\so2, $\\HH$S and HCS\\p\\ from some of the diffuse clouds which occult our well-studied sample of compact extragalactic mm-wave continuum sources.  Our observations of SO, $\\HH$S and HCS\\p\\ represent the first detections of these species in diffuse clouds; \\so2\\ was not detected at all. We find a typical value N(CS) $\\approx 0.5 - 2.0 \\times 10^{12}~\\pcc$ which is actually much smaller than the values derived previously, along very different lines of sight, from interpretation of CS J=2-1 emission lines.  CS forms somewhat sluggishly and is occasionally absent even in features with appreciable N(\\hcop).  But for lines of sight where CS is found, N(CS)/N(\\hcop) $\\approx 2\\pm1$ or X(CS) $\\approx 4 \\times10^{-9}$.  N(CS) correlates well and varies about linearly with N(HCN) (N(CS)/N(HCN) = $0.7\\pm0.3$) though, again, lines of sight with appreciable N(HCN) occasionally lack CS.  CS correlates well with $\\HH$S (N(CS)/N($\\HH$S) $ = 6\\pm1$) and marginally with SO (N(CS)/N(SO) = $1.7\\pm0.8$).  For the two high column density features observed toward 3C111 we find N(CS)/N(HCS$^+$) = 13.3$\\pm1.0$. CS can easily be shown to form from the observed amounts of HCS\\p\\ $via$ electron recombination in cool, quiescent gas but the obvious gas-phase routes to formation of HCS\\p\\ fail by factors of 25 or more. ",
        "introduction": " ",
        "conclusions": "\\begin{figure} \\psfig{figure=ms2095f2.eps,height=9.4cm} \\caption[]{Digest of detected SO absorption features.} \\end{figure}  Observationally, Snow's $Copernicus$ limits on SH toward $o$ Per, N(SH) $< 10^{12}~\\pcc$ \\citep{Sno75}, and on CS toward \\zoph, N(CS) $< 2.6 \\times 10^{13}~\\pcc$ \\citep{Sno76}, constituted the entire body of knowledge of sulfur chemistry in diffuse clouds for some time.  \\cite{FerRou+86} then showed that N(CS\\p) $< 1.9\\times10^{11}~\\pcc$ toward \\zoph, amending an earlier, tentative detection.  Eventually, \\cite{DrdKna+89} detected very weak (0.01 - 0.10 K) CS J=2-1 emission toward several objects occulted by supposedly diffuse gas (A$_V \\la 1$ mag), including $o$ Per and HD210121 (T$_{\\rm A}^\\ast = 0.03-0.04$ K, N(CS) $ =  4-100 \\times 10^{12}~\\pcc$, N(CS)/N($\\HH$) $ \\ge 0.4 - 10 \\times 10^{-8}$), and (somewhat tentatively) \\zoph\\ (\\Ta = 0.01 K, N(CS) $ =  0.7-5 \\times 10^{12}~\\pcc$, N(CS)/N($\\HH$) $ = 0.15 - 1.1 \\times 10^{-8}$).  The large ranges in derived CS column density in the work of \\cite{DrdKna+89} reflect a sensitivity to the assumed physical conditions, but the implied relative abundances are really very high -- in fact, quite comparable to or even larger than those found in dark clouds and GMC's. \\footnote{As is evident from our work, where the typical CS column density is N(CS) $= 1-2 \\times 10^{12} ~\\pcc$,  the higher CS abundances found by \\cite{DrdKna+89} suggest that either the density was underestimated or that the gas sampled in emission is considerably less diffuse than was believed at the time.} Nonetheless, \\cite{DrdKna+89} concluded that CS could be made in the required amount by ordinary gas-phase ion-molecule chemistry acting in quiescent diffuse/translucent gas with n(H) $= 200-1000~\\pccc$ and $\\Tk = 30-50$ K, and they predicted (for the line of sight toward \\zoph) the diffuse-cloud abundances of some of the other sulfur-bearing species known to exist in molecular clouds, {\\it i.e.} N(CS)/N(HCS$^+$) = 40 - 80, N(CS)/N(SO) = 2800 - 3800, N(CS)/N($\\HH$S) $\\approx$ N(CS)/N(SO$_2$) $\\approx 10^9$. For these ratios, the observed values from the present data (along other directions) are 13, 2, 6, and $>$ 2.  \\subsection{CS emission toward \\zoph}  \\begin{figure} \\psfig{figure=ms2095f3.eps,height=9.4cm} \\caption[]{Digest of detected $\\HH$S and (lower right) HCS\\p\\ absorption features.  The HCS\\p\\ profile has been scaled upward by a factor 3.} \\end{figure}  We tried and failed (see \\cite{Lis97}) to confirm the interesting but somewhat tentative detection of $0.013\\pm0.003$ K CS J=2-1 emission at the somewhat unusual velocity of -1.29 \\kms\\ toward \\zoph\\ by \\cite{DrdKna+89}.  Our spectra of \\hcop\\ and CS at two positions are shown in Fig. 6.  At the \\hcop\\ emission peak 30\\arcmin\\ South of the star, CS emission is at least an order of magnitude weaker than that of \\hcop. If this proportionality were preserved, CS emission toward the star would be six times weaker than claimed by \\cite{DrdKna+89}. Toward the star, the $1\\sigma$ rms fluctuation in integrated CS line intensity in our data is 0.005 K \\kms, but there is no signal. So the $4\\sigma$ integrated intensity detection quoted by \\cite{DrdKna+89}, I$_{\\rm CS} = 0.020 \\pm0.005$ K \\kms, is a $4\\sigma$ upper limit in our data.    Table 3 presents a comparison of relative abundances between the diffuse gas observed here and representative results for the well-studied dark clouds TMC-1 and L134N.  It should be noted that there are large variations in the CS/SO ratio between and within individual dense sources observed in emission (including these two dark clouds) and that TMC-1 and L134N, whose CS/SO ratios differ by about a factor 40, represent the extremes of the observed range \\citep{NilHja+00,GerFal+97}.  The abundances of CS and \\hcop\\ relative to $\\HH$ in diffuse clouds are about 25\\%-50\\% that inferred for TMC-1.  The abundance ratios relative to \\hcop\\ are essentially identical for our sample of diffuse clouds and the representative conditions quoted for TMC-1, with the possible exception of an overabundance of $\\HH$S in our objects. For the dense clouds and GMC studied by \\cite{NilHja+00}, the typical SO/CS ratio is 0.2 (i.e. it is TMC-1-like) with only one-fifth at or above unity; this is just what is seen in diffuse clouds as well.  But for translucent and dark clouds, ratios SO/CS $<$ 1 are actually the exception as detailed in Fig. 3c of \\cite{GerFal+97} which includes the data collected by \\cite{Tur95,Tur96}.  \\begin{table} \\caption[]{Elemental and Molecular abundance ratios} { \\begin{tabular}{lcccc} \\hline Ratio &Solar& \\zoph & Species & Diffuse  \\\\ \\hline C/O    &  0.50   &  0.50   & CS/SO    & $1.7\\pm0.8$  \\\\ &         &        & CH/OH    & 0.48,0.52,0.24  \\\\ C/N    &  3.80   &  1.74  & CH/NH    & 22,28  \\\\ &         &        & C$_2$/CN &  9  \\\\ &         &        & \\cch/HCN & 5   \\\\ S/O    &  0.025  &  0.10   & HCS\\p/\\hcop\\ & 0.055,0.070  \\\\ S/N    &  0.20   &  0.36  & CS/CN &$0.1\\pm0.05$   \\\\ \\hline \\end{tabular}} \\\\ \\zoph\\ elemental abundances from Savage and Sembach (1996) \\\\ OH and CH abundances from \\cite{BlaVan86} \\\\ NH from \\cite{CraWil97} and \\cite{MeyRot91} \\\\ C$_2$/CN from \\cite{FedStr+94}, see Paper II, Fig. 3 \\\\ \\end{table}  It is now possible to make comparisons of the abundances of several species which differ in structure by the replacement of one atom with that of a different element, as for example HCS\\p\\ and HCO\\p, or CS and SO, {\\it etc.}.  Table 4 points out that these abundance ratios are not very different from that of the elements which are exchanged, the outstanding exception being the CH/NH ratio which is substantially greater than either C$_2$/CN  or \\cch/HCN."
    },
    "0201/astro-ph0201480_arXiv.txt": {
        "abstract": "We present the first orbital elements for the massive close binary, HD~115071, a double-lined spectroscopic binary in a circular orbit with a period of $2.73135 \\pm 0.00003$ days.  The orbital semiamplitudes indicate a mass ratio of $M_2/M_1 = 0.58 \\pm 0.02$ and yet the stars have similar luminosities.   We used a Doppler tomography algorithm to reconstruct the individual component optical spectra, and we applied well known criteria to arrive at classifications of O9.5~V and B0.2~III for the primary and secondary, respectively.   We present models of the {\\it Hipparcos} light curve of the ellipsoidal variations caused by the tidal distortion of the secondary, and the best fit model for a Roche-filling secondary occurs for an inclination of $i=48\\fdg7 \\pm 2\\fdg1$. The resulting masses are  $11.6\\pm1.1 M_\\odot$ and $6.7\\pm 0.7 M_\\odot$ for the primary and secondary, respectively, so that both stars are very overluminous for their mass. The system is one of only a few known semi-detached, Algol-type binaries that contain O-stars.  We suggest that the binary has recently emerged from extensive mass transfer (possibly through a delayed contact and common envelope process). ",
        "introduction": "%  The hot, massive star, HD~115071 (V961~Cen, LS~2998, HIP~64737), is found in the sky close to the open cluster, Stock~16 \\citep{tur85}, and is classified as O9.5~V by \\citet{hou75} and B0.5~Vn by \\citet{gar77}.  The star is not a known visual binary \\citep{mas98} but early measurements by spectroscopists indicated it is radial velocity variable and a probable spectroscopic binary \\citep{cru74,con77}. The proof of its binary nature came relatively recently in studies by \\citet{pen96} and \\citet{how97}.  Both papers presented a cross-correlation analysis of a single, high dispersion, UV spectrum made with the {\\it International Ultraviolet Explorer Satellite} ({\\it IUE}) that demonstrated that the system is in fact a double-lined binary.   \\citet{sti01} measured the radial velocities of the components in this spectrum and proposed an orbital period of 2.73126~d based upon a light curve constructed from {\\it Hipparcos} photometry. \\citet{llo01} present a model of the light curve, and they argue that the system has evolved through Case~A mass transfer (commencing during core H burning of the donor star).  The details and outcomes of Roche lobe overflow (RLOF) in massive binaries are still subjects of considerable debate \\citep{wel01}, and thus, the orbital and physical parameters of a system like HD~115071 are of great interest.  Here we present the first double-lined orbital solution for the binary (\\S3) based upon new high quality optical spectra.  We apply a version of the Doppler tomography algorithm (which we have used to good effect with UV spectra in prior papers in this series) to reconstruct the individual spectra of both components, from which we determine the spectral classifications, projected rotational velocities, and flux ratio (\\S4). We also present a light curve analysis constrained by the spectroscopic results that allows us to estimate the stellar masses (\\S5).  These masses are much lower than expected, and we discuss the evolutionary implications in \\S6. ",
        "conclusions": "%  The first striking result from our analysis is the very low mass we find for both components.  The stars have temperatures and luminosities that are associated with masses of 18 and $15 M_\\odot$ for the primary and secondary, respectively, in the single star evolutionary tracks calculated by \\citet{sch92}. (These estimates would be slightly reduced using evolutionary models that include rotation; \\citet{heg00}, \\citet{mey00}.) The secondary, in particular, has a luminosity characteristic of a star more than twice as massive than we find (Table 3).  The second remarkable fact is that the secondary star has a spectral classification indicating it has evolved away from the main sequence.   Thus, HD~115071 presents the classical ``Algol paradox'' that the lower mass component is the more evolved one, and we suggest the same solution of the paradox holds here as well, i.e., that the evolved component was originally the more massive object, but suffered significant mass transfer to its neighbor.  There are only a small number of O-stars that are known to be members of interacting binaries, and we compare in Table~4 the properties of the components in HD~115071 with those of the four other known semi-detached binaries that contain O-type stars \\citep{hil87,hh98}.  We excluded from this list contact or over-contact systems and those binaries in which both components are evolved \\citep{van98}.  All the systems in Table~4 share a number of common properties: the mass donor appears as an evolved star, the donor star is overluminous for its mass, the donor fills its Roche volume, and the mass gainer is a late O-type, main sequence star. It is remarkable that all the donor stars have comparable luminosity, $\\log L/L_\\odot \\approx 4.5$, despite their wide range in mass and radius.   Evolutionary models generally predict that the post-RLOF luminosity of the donor is comparable to its zero age main sequence (ZAMS) luminosity \\citep{van98,wel01}, and so these donors probably began life as B0~V stars with masses in the range 14 -- $20 M_\\odot$.  Since the donors were originally the more massive component, the gainers were probably also B-type stars that were promoted to their current O-type status through mass transfer. It is also curious that no semi-detached systems are known with primaries earlier than type O8~V.  Either this stage is extremely rapid in more massive systems or the donor stars take on a different appearance than they do in Algol-type systems (perhaps as a O-star plus Wolf-Rayet star binary; \\citet{van98}).  \\placetable{tab4}      %  Evolutionary models give us some guidance about the initial masses in HD~115071.  \\citet{del94} give a relationship between the final, post-RLOF mass and the initial ZAMS mass, and this yields an estimate of $14.8 M_\\odot$ for the initial mass of the donor star.  If we further assume that 50\\% of the donor's mass loss was accreted by gainer and the rest lost from the system \\citep{meu89,del94}, then the original total mass was $22.4 M_\\odot$ and the original gainer mass was $7.6 M_\\odot$. Thus, the system probably began with a relatively low mass ratio, $M_g/M_d \\approx 0.5$.  The theoretical models of binary evolution by \\citet{wel01} offer some guidance in the interpretation of our results. \\citet{wel01} describe the evolution of several very close systems that begin RLOF during core H-burning (Case A).   Their models suggest that a mass reversal similar to what we find in HD~115071 can occur in Case A, but the resulting systems generally have a much wider orbit and more extreme mass ratio. Another possibility is that the system began RLOF after completion of core H burning (Case B).  The initial period would have been much larger, but the system then shrunk to its current dimensions during a common envelope phase in which the donor's envelope would have been ejected from the system.  This would explain the current mass and luminosity of the donor star, but it does not account for the huge overluminosity of the contemporary primary star, which is the most overluminous star of any of the gainers in Table 4.   \\citet{wel01} point out one other hybrid scheme they call ``delayed contact'' in which mass transfer begins conservatively until the donor develops a convective envelope and the binary enters the common envelope stage. This scenario would explain the observed low mass of the donor and the short orbital period, and the overluminosity of the gainer would result from compression and/or mixing related to mass accretion.  The best fit of the light curve suggests that the secondary donor star is Roche-filling, and so the system may still be experiencing active mass transfer.  Observations of any H$\\alpha$ emission \\citep{tha97} or IR excess \\citep{geh95} would provide valuable clues about the mass loss and/or mass transfer processes that might be occurring presently in this exceptional binary system."
    },
    "0201/astro-ph0201449_arXiv.txt": {
        "abstract": "{The detection of 22\\,GHz water vapor emission from IC~342 is reported, raising the detection rate among northern galaxies with IRAS point source fluxes $S_{\\rm 100\\mu m}$ $>$ 50\\,Jy to 16\\%. The maser, associated with a star forming region $\\sim$10--15\\arcsec\\ west of the nucleus, consists of a single 0.5\\,km\\,s$^{-1}$ wide feature and reaches an isotropic luminosity of 10$^{-2}$\\,L$_{\\odot}$ ($D$=1.8\\,Mpc). If the time variability is intrinsic, the maser size is $\\la1.5\\times10^{16}$\\,cm ($\\la$0.5\\,mas) which corresponds to a brightness temperature of $\\ga$10$^{9}$\\,K. The linewidth, luminosity, and rapid variability are reminiscent of the 8\\,km\\,s$^{-1}$ super maser in Orion-KL. A velocity shift of 1\\,km\\,s$^{-1}$ within two weeks and subsequent rapid fading is explained in terms of a chance alignment of two dense molecular clouds. Observations at 22 GHz toward Maffei~2 are also reported, yielding a 5$\\sigma$ upper limit of 25\\,mJy for a channel spacing of 1.05\\,km\\,s$^{-1}$. ",
        "introduction": "Isotropic luminosities of extragalactic 22\\,GHz $\\rm H_{2}O$ masers span a huge range (see e.g. Ho et al.\\ \\cite{ho87}; Koekemoer et al.\\ \\cite{koeke95}; Braatz et al.\\ \\cite{braatz96}), from $\\rm 10^{-2} - 10$~$\\rm L_{\\odot}$ (the `kilomasers' consistent with the most intense masers seen in our Galaxy) to 10 -- 500~$\\rm L_{\\odot}$ (the `megamasers', found in LINERs and Seyfert 2s) and up to 6000 $\\rm L_{\\odot}$ (the `gigamaser' in TXS\\,2226-184). Interferometric observations, when available, show that the most luminous $\\rm H_{2}O$ sources arise from the nuclear region of their host galaxies, from circumnuclear tori, from entrained material near the surface of the nuclear jet, or by ambient gas amplifying the continuum emission of the jet. The weaker masers often mark locations of high mass star formation (for M~33, IC~10, and M~82, see e.g. Huchtmeier et al.\\ \\cite{hucht78}; Henkel et al.\\ \\cite{henkel86}; Baudry \\& Brouillet \\cite{baundry96}). The nature of the kilomasers in NGC~253 and M~51 (see Ho et al.\\ \\cite{ho87}; Nakai \\& Kasuga \\cite{nakai88}) remained an enigma for more than a decade. Recently, however, it has been shown that the maser in M~51 is of nuclear origin (Hagiwara et al.\\ \\cite{hagi01}), suggesting that there might be a family of as yet unexplored `weak' H$_2$O sources located near the nuclear engine of Seyfert 2 and LINER galaxies.  Extragalactic H$_2$O masers are preferentially detected in nearby galaxies that are bright in the infrared (Braatz et al.\\ \\cite{braatz97}). While nuclear masers are of obvious interest, non-nuclear kilomasers are also important for a number of reasons: these sources allow us to pinpoint sites of massive star formation, to measure the velocity vectors of these regions through VLBI proper motion studies, and to determine true distances through complementary measurements of proper motion and radial velocity (e.g. Greenhill et al.\\ \\cite{greenhill93}). We have therefore observed the nearby spiral galaxies IC~342 and Maffei~2, both of which exhibit prominent nuclear bars and strong molecular, infrared, and radio continuum emission (e.g. Hurt \\& Turner\\ \\cite{hurt91}; Hurt et al.\\ \\cite{hurt93}; Turner \\& Ho\\ \\cite{turner94}; Henkel et al.\\ \\cite{henkel00}; Meier et al.\\ \\cite{meier00}; Meier \\& Turner\\ \\cite{meier01}; Schulz et al.\\ \\cite{schulz01}). In the following we report the results of our observations. ",
        "conclusions": "Having detected water vapor emission in the galaxy IC~342, we have obtained the following main results:  \\begin{enumerate}  \\item The maser arises from a location 10--15\\arcsec\\ to the west of the center of the galaxy. It is thus a maser associated with a powerful star forming region at a projected distance of $\\sim$100\\,pc ($D$ = 1.8\\,Mpc) from the nucleus.  \\item Luminosity (10$^{-2}$\\,L$_{\\odot}$), linewidth (0.5\\,km\\,s$^{-1}$), and rapid flux density variations are reminiscent of the flaring 8\\,km\\,s$^{-1}$ super maser feature in Orion-KL. It seems that we have observed an exceptional outburst of a maser that is usually well below the detection threshhold.  \\item Time variability, if intrinsic, yields a maser size of $\\la 1.5\\times10^{16}$\\,cm$^{-3}$ ($\\la$0.5\\,mas), and a brightness temperature of $\\ga$10$^{9}$\\,K.  \\item A velocity shift of 1\\,km\\,s$^{-1}$ within 16 days and a subsequent decrease in flux density by at least a factor of three within 5 days can be explained by a chance alignment of two dense molecular clouds along the line-of-sight, both of them associated with IC~342.  \\item Among galaxies with IRAS point source flux densities of $S_{\\rm 100\\mu m}$ $>$ 50\\,Jy, 16\\% are now known to contain masers in their inner regions. Since deep H$_2$O integrations have been obtained towards only a few of them, more detections from this sample can be expected in the near future.  \\end{enumerate}"
    },
    "0201/astro-ph0201213_arXiv.txt": {
        "abstract": "Radial-velocity measurements and sine-curve fits to the orbital velocity variations are presented for the sixth set of ten close binary systems: SV~Cam, EE~Cet, KR~Com, V410~Cyg, GM~Dra, V972~Her, ET~Leo, FS~Leo, V2388~Oph, II~UMa. All systems except FS~Leo are double-lined spectroscopic binaries. The type of FS~Leo is unknown while SV~Cam is a close, detached binary; all remaining systems are contact binaries. Eight binaries (all except SV~Cam and V401~Cyg) are the recent photometric discoveries of the Hipparcos satellite project. Five systems, EE~Cet, KR~Com, V401~Cyg, V2388~Oph, II~UMa, are members of visual/spectroscopic triple systems. We were able to observe EE~Cet separate from its companion, but in the remaining four triple systems we could separate the spectral components only through the use of the broadening-function approach. Several of the studied systems are prime candidates for combined light and radial-velocity synthesis solutions. ",
        "introduction": "\\label{sec1}  This paper is a continuation in a series of papers of radial-velocity studies of close binary stars: \\citet[Paper I]{ddo1}, \\citet[Paper II]{ddo2} and \\citet[Paper III]{ddo3}, \\citet[Paper IV]{ddo4}, \\citet[Paper V]{ddo5}. The main goals and motivations are described in these papers. The companion paper of \\citet[Paper VII]{ddo7} describes the technical details and methods of data reductions as well as the resulting measurement uncertainties, based on data in the six previous papers.  This paper is structured in the same way as Papers~I -- V in that most of the data for the observed binaries are in two tables consisting of the radial-velocity measurements (Table~\\ref{tab1}) and their sine-curve solutions (Table~\\ref{tab2}). Section~\\ref{sec2} of the paper contains brief summaries of previous studies for individual systems. Figures~\\ref{fig1} -- \\ref{fig3} show the data and the radial velocity solutions.  The observations reported in this paper have been collected between October 1997 and June 2001; the ranges of dates for individual systems can be found in Table~\\ref{tab1}. All systems discussed in this paper, except SV~Cen, have been observed for radial-velocity variations for the first time. We have derived the radial velocities in the same way as described in previous papers. See Paper~IV and Paper~VII for discussion of the broadening-function approach used in the derivation of the radial-velocity orbit parameters: the amplitudes, $K_i$, the center-of-mass velocity, $V_0$, and the time-of-eclipse epoch, $T_0$.  The data in Table~\\ref{tab2} are organized in the same manner as in previous papers. In addition to the parameters of spectroscopic orbits, the table provides information about the relation between the spectroscopically observed epoch of the primary-eclipse T$_0$ and the recent photometric determinations in the form of the (O--C) deviations for the number of elapsed periods E. It also contains our new spectral classifications of the program objects.  For further technical details and conventions used in the paper, please refer to Papers I -- V and VII of this series.  \\placetable{tab1} \\placetable{tab2}  \\placefigure{fig1} \\placefigure{fig2} \\placefigure{fig3} ",
        "conclusions": "The paper presents radial-velocity data and orbital solutions for the sixth group of ten close binary systems that we observed at the David Dunlap Observatory. Only the detached, short-period system SV~Cam was observed spectroscopically before. All systems except FS~Leo are double-lined (SB2) binaries with visible spectral lines of both components and all, with the additional exception of SV~Cam, are contact binaries.  Although our selection of our targets is quite unsystematic and is driven only by the brightness above the limit of about 11 mag.\\ and the shortness of the orbital period (below one day), we continue seeing very interesting objects among the bright, photometrically discovered binaries. Eight of the systems are new photometric discoveries of the Hipparcos project. The magnitude-limited nature of the HIP survey has led to an emphasis on relatively luminous, massive, early-type (middle-A to early-F) contact systems. Because the Hipparcos mission discovered mostly small photometric amplitude systems, which were overlooked in previous whole-sky surveys, we tend to observe systems somewhat pre-selected in that they either have low orbital inclinations, such as V972~Her or ET~Leo, or are members of triple systems as for KR~Com, V401~Cyg, V2388~Oph and II~UMa. In the latter case, while radial velocity amplitudes are large, the diminished photometric variation is due to the ``dilution'' of the variability signal in the combined light of the triple system. In addition to these four triple systems, we also observed separately a component in a wider visual binary, EE~Cet; the other component of this system appears to be radial-velocity variable whose properties remain to be studied."
    },
    "0201/astro-ph0201025_arXiv.txt": {
        "abstract": "Using the sunspot numbers reported during the Maunder minimum and the empirical relations between the mode frequencies and solar activity indices, the variations in the total solar irradiance and 10.7 cm radio flux  for the period 1645 to 1715 is estimated.  We find that the total solar irradiance and radio flux during the Maunder minimum decreased by 0.19\\% and  52\\% respectively, as compared to the values for solar cycle 22. ",
        "introduction": "The period of low solar activity  between 1645 to 1715 is commonly known as the Maunder minimum and has attracted attention of researchers as it put forward many unexpected questions. During this period,  the sunspots were infrequent and as a result it was assumed that the solar dynamo and the solar wind were either weak or switched off. Anomalies in the surface differential rotation during the Maunder minimum \\cite{ribes93} also point out the global changes in the internal solar dynamo mechanism. Measurements of $^{10}$Be concentration in the Dye 3 ice core \\cite{beer98} recently showed that the magnetic cycles persisted throughout the Maunder minimum, although the Sun's over all activity was drastically reduced. This period is also associated with the so called ``little ice age'' in Europe and \\inlinecite{eddy76} suggested that it might be due to a decrease in solar irradiance. Many authors have estimated the total solar irradiance during the Maunder minimum and find a decrease in the range of 0.1\\% at a time of relatively high activity down to 1\\% at a time of no or low activity.  In a different approach, we estimate the change in the total solar irradiance and 10.7 cm radio flux during the Maunder minimum. This is achieved by calculating the  {\\it p}-mode frequencies using the historic record of sunspots and the empirical relations derived by \\inlinecite{jain99}.  Since these relations between mode frequencies and solar activity indices are shown to be independent of solar cycles, these can be reliably used to estimate variations in solar irradiance and radio flux.  \\begin{figure} \\psfig{file=fig1.eps,height=216pt}% \\caption{Estimated frequency shift for the period 1610--1750 (solid line) and 1965--1998 (dashed line) with reference to the annual mean of 1996, using  the mean annual sunspot number from Eddy (1976) and Equation~(1). }  \\label{fig1} \\end{figure} ",
        "conclusions": "Using the  historical sunspot data and the  empirical relations between the {\\it p}-mode frequencies and activity indices derived by \\inlinecite{jain99}, it is estimated that during the Maunder minimum from 1645 to 1715 AD, the mode frequencies showed a maximum change  of 43 $\\pm$ 3 nHz. Using the derived mode frequencies as a proxy of the solar activity indicators, we have further estimated the total solar irradiance and 10.7 cm radio flux and find that these are lowered by 0.19\\% and 52\\% respectively than the average values of the solar cycle 22. The decrease in {\\it p}-mode frequency and irradiance during Maunder minimum probably indicates stability in the dynamical processes occuring in the solar convection zone and even deep down up to the core."
    },
    "0201/astro-ph0201031_arXiv.txt": {
        "abstract": "\\noindent The SHARP/NTT stellar proper motion data now cover an interval from 1992 to 2000 and allow us to determine orbital accelerations for some of the most central stars. We confirm the stellar acceleration measurements obtained by Ghez et al. (2000) with NIRC at the Keck telescope. Our analysis differs in 3 main points from that of Ghez et al.: 1) We combine the high precision but shorter time scale NIRC/Keck data with the lower precision but longer time scale SHARP/NTT data set; 2) We statistically correct the observed accelerations for geometrical projection effects; 3) We exclude star S8 from the analysis of the amount and position of the central mass. \\\\ \\\\ From the combined SHARP/NTT and NIRC/Keck data sets we show that the stars S2, and most likely  S1 and S8 as well, are on bound, fairly inclined ($60^o<i<80^o$), and eccentric ($0.4<e<0.95$) orbits around a central dark mass. The combination of both data sets results in a position of this central mass of 48$^{+54}_{-24}$~$mas$ E and 18$^{+42}_{-61}$~$mas$ S of the nominal radio position of Sgr~A*. The mean statistically corrected enclosed mass derived from accelerations is $M_{acc}=(5 \\pm 3) \\times 10^6$\\solm with current radial separations of S1 and S2 from SgrA* of about $8-10~mpc$. This enclosed mass estimate is derived from {\\it individual} stellar orbits as close to the massive black hole at the center of the Milky Way as currently possible. Although the uncertainties are large this estimate is fully consistent with the enclosed mass range of $(2.6-3.3)\\times10^6$\\solm~derived by Genzel et al. (2000) from radial and/or proper motion velocities of a homogenized sample of sources. \\\\ \\\\ Star S8 was excluded from the analysis, since for the current proper motion velocity and radial separation from the center we find that the measured acceleration requires orbital motion around a compact object with a mass in excess of $3\\times10^6$\\solm. The data suggest that this star {\\it either} was or is subject to a close interaction with a different object {\\it or} that its position measurements are influenced by the emission of a different cluster star. Therefore, we base the analysis of the enclosed mass solely on the available data for S1 and S2. We also discuss two late type stars with projected separations from SgrA* of about 0.5'' and 1''. In addition to proper motions these stars have known radial velocities. Orbit calculations indicate that those stars are very likely at larger physical distances from the center and part of the larger scale central stellar cluster with a core radius of approximately $0.3~pc$. ",
        "introduction": "INTRODUCTION} \\normalsize High resolution near-infrared imaging and spectroscopy with large telescopes resulted in a determination of the amount and concentration of mass at the center of the Milky Way (Sellgren et al. 1990, Krabbe et al., 1995, Haller et al. 1996, Eckart \\& Genzel 1996, Genzel et al. 1998, Ghez et al. 1998, Genzel et al. 2000). Diffraction limited images, proper motions and most recently the detection of acceleration of stars in the vicinity of the compact radio source Sgr~A$^*$ (Ghez et al. 2000, Eckart et al. 2000a) have led to the conclusion that this source is associated with a central black hole with a mass of about 3$\\times$10$^{{ 6}}$ \\solm. Using the MPE speckle camera SHARP at the 3.5 m New Technology Telescope (NTT) of the European Southern Observatory (ESO) from 1992 to 2000 we have been conducting a program to study the properties of the central nuclear stellar cluster via near-infrared high spatial resolution measurements. This program has resulted in the very first detection of proper motions of stars that correspond to velocities of up to 1400~km/s in the central arcsecond in the vicinity of Sgr~A$^*$ (Eckart \\& Genzel 1996, 1997). These results had been confirmed by Ghez et al. (1998) using NIRC on the Keck telescope. On the 1.6$\\sigma$ to 3$\\sigma$ level we have now detected orbital curvatures of the stars S1, S2, and S8 which confirm the recent results by Ghez et al. (2000). \\\\ \\\\ Over the past 8 years the observed motions translate into a two dimensional velocity dispersion of the stars in the central arcsecond (corrected for the measurement error) of the order of 600 km/s. Overall the stellar motions do not deviate strongly from isotropy and are consistent with a spherical isothermal stellar cluster (Genzel et al. 2000). However, a small deviation from isotropy is found for the sky-projected velocity components of the young, early type stars. Most of the bright HeI emission line stars are on tangential orbits. This overall rotation could be a remnant of the original angular momentum pattern in the interstellar cloud from which these stars were formed. The projected radial and tangential proper motions of the fainter, fast moving stars within $\\approx${}1'' from SgrA* (the `SgrA* cluster') suggest that they may largely be moving on very eccentric  orbits. Speckle spectroscopy with SHARP at the NTT (Genzel et al. 1997) and slit spectroscopy with ISAAC at the VLT (Eckart, Ott, Genzel 1999, Figer et al. 2000) suggest that several of them are early type stars. This is consistent with the idea that these stars are members of the early type cluster with small angular momentum and therefore fell into the immediate vicinity of SgrA* (Genzel et al. 2000, Gerhard 2000, see also discussion in Ghez et al. 2000).  The detection of orbital curvature for 3 central stars now allows us for the first time to constrain their orbital elements and to investigate the central mass distribution at the smallest currently accessible separations from Sgr~A$^*$. In the following subsection \\ref{sec1.1} we give a summary of why it is worthwhile to determine stellar orbits near the dark mass at the Galactic Center and give a general description of how that can be done best. In section \\ref{sec2} we present the new proper motion and acceleration data and derive the position of the central dark mass. In section \\ref{sec3} we show how the acceleration measurements can be used to obtain an enclosed mass estimate. In section \\ref{sec4} we then make use of the enclosed mass and its position and discuss possible orbital solutions for several of the central stars. Finally a summary and conclusions are given in section \\ref{sec5}.  \\subsection{ \\label{sec1.1} Stellar orbits and the central mass distribution} \\normalsize  The point source sensitivities currently reached in near infrared observations of the central stellar cluster are of the order of $K\\sim16-17$ at an angular resolution of 50 - 150 mas (Eckart \\& Genzel 1996, Ghez et al. 1998, Genzel et al. 2000, Ghez et al. 2000). These current limits are mainly a result of the imaging techniques at the diffraction limit and not due to the sensitivity of the instrumentation. With these values the analysis of the central mass distribution for separations of less than 0.5'' ($\\sim$20~mpc) from the radio point source SgrA* has to be based on detailed measurements of individual stellar orbits. We assume that the trajectories of stars can be approximated by Keplerian shaped orbits around a central mass.  Relative Keplerian orbits are described by 6 orbital elements: three angles, the eccentricity $e$, the semi-major axis $a$ and a reference time $T$ for a defined position (i.e. the periastron). Currently for a few of the early type high velocity S-sources as well as two late type stars that are close to the center in projection, our detailed orbital analysis (section \\ref{sec4}) can only be based on 5 known quantities that constrain the orbits. In addition to the projected positions and proper motions these are the orbital curvatures (for the S-sources) and the line of sight velocities (for the late type stars). However, for visual binaries - as is the case for a combination of the mass at the center of the Galaxy and any of the fast moving stars in its vicinity - all these orbital elements can be derived from a sufficiently large section of the orbit, except for the sign of the inclination. This has to be inferred from Doppler shifts via spectroscopy. The visual binaries situation holds if the exact position of the black hole on the near-infrared images is known (see Menten et al. 1996). The orbital period and central mass can then be determined via Kepler's third law and the equation of motion i.e. from Newton's law. Here we present a first step to reach that goal using the combined SHARP/NTT and NIRC/Keck data sets.  The orbital analysis in the presence of a possibly partially extended central mass has to await the detection of fainter stars ($K\\ge 17$) closer ($<2~mpc$) to SgrA* or the availability of long period measurements (of the order of 100 years or more; see section \\ref{sec4.2}) of the presently known high velocity stars within the central arcsecond. The expected orbits have been discussed in the context of the Galactic Center by Rubilar \\& Eckart (2001, see also Fragile \\& Mathews, 2000, Munyaneza, Tsiklauri \\& Viollier, 1998, and Jaroszy\\'nski, 1998). The shortest orbital time scales that could be observed in the near future will be of the order of only a few months. These measurements are well within reach using the new instrumentation (VLT, Keck and LBT interferometers; see Rubilar \\& Eckart 2001, Eckart, Ott, Genzel 2001a, 2001b). ",
        "conclusions": ""
    },
    "0201/astro-ph0201207_arXiv.txt": {
        "abstract": "We derive the number density evolution of massive field galaxies in the redshift range $0.4 < z < 1.2$ using the $K$-band selected field galaxy sample from the Munich Near-IR Cluster Survey (MUNICS).  We rely on spectroscopically calibrated photometric redshifts to determine distances and absolute magnitudes in the rest-frame $K$-band. To assign mass-to-light ratios, we use two different approaches.  First, we use an approach which maximizes the stellar mass for any $K$-band luminosity at any redshift. We take the mass-to-light ratio, \\MLK , of a Simple Stellar Population (SSP) which is as old as the universe at the galaxy's redshift as a likely upper limit.  Second, we assign each galaxy a mass-to-light ratio by fitting the galaxy's colours against a grid of composite stellar population models and taking their \\MLK.  We compute the number density of galaxies more massive than $2\\times 10^{10} h^{-2} \\Msun$, $5\\times 10^{10} h^{-2} \\Msun$, and $1\\times 10^{11} h^{-2} \\Msun$, finding that the integrated stellar mass function is roughly constant for the lowest mass limit and that it decreases with redshift by a factor of $\\sim 3$ and by a factor of $\\sim 6$ for the two higher mass limits, respectively. This finding is in qualitative agreement with models of hierarchical galaxy formation, which predict that the number density of $\\sim M^{*}$ objects is fairly constant while it decreases faster for more massive systems over the redshift range our data probe. ",
        "introduction": "The traditional observables used to characterise galaxies are unsuitable for studying the assembly history of galaxies, one of the most fundamental predictions of CDM models, since these observables may be transient. The best observable for this aim is, in principle, total mass, which is on the other hand very hard to measure. It has been argued that the best available surrogate accessible to direct observation is the near-IR $K$-band luminosity of a galaxy which reflects the mass of the underlying stellar population and is least sensitive to bursts of star formation and dust extinction \\citep{RR93,KC98a,BE00}. The main uncertainty involved in the conversion of $K$-band light to mass is due to the age of the population, amounting to only a factor of two in mass uncertainty for populations older than $\\sim 3$~Gyr.  The galaxy sample used here is a subsample of the MUNICS survey, selected for best photometric homogeneity, good seeing, and similar depth.  Furthermore, in each of the remaining survey patches, areas close to the image borders in any passband, areas around bright stars, and regions suffering from blooming are excluded. The subsample covers 0.27 square degrees in $V$ (23.5), $R$ (23.5), $I$ (22.5), $J$ (21.5), and $K$ (19.5); the magnitudes are in the Vega system and refer to 50\\% completeness for point sources.  The final catalog covers an area of 997.7 square arc minutes and contains 5132 galaxies. The fields included in this analysis are S2, S3f5--8, S5, S6, and S7f5--8. See Table 1 in \\cite{MUNICS1} for nomenclature and further information on the survey fields.  The distances to the galaxies are derived using spectroscopically calibrated photometric redshifts. A comparison of spectroscopic and photometric redshifts is shown in \\citet{MUNICS3}.  We derive stellar masses by converting rest-frame $K$-band luminosities to mass using two different approaches to model the mass-to-light ratios of the galaxies.  We discuss the resulting integrated stellar mass functions at different mass limits and their evolution with redshift.  We assume $\\Omega_M = 0.3$, $\\Omega_{\\Lambda} = 0.7$ throughout this work. We write Hubble's Constant as $H_0 = 100\\ h\\ \\mathrm{km\\ s^{-1}\\ Mpc^{-1}}$, using $h = 0.65$ unless the quantities in question can be written in a form explicitly depending on $h$. ",
        "conclusions": ""
    },
    "0201/astro-ph0201494_arXiv.txt": {
        "abstract": "We present the results of ultraviolet echelle spectroscopy of a sample of fifty nine F,G,K and M stars from the Anglo-Australian Planet Search target list. Ca\\,\\II\\ activity indices, which are essential in the interpretation of planet detection claims, have been determined for these stars and placed on the Mount Wilson \\rhk\\ system.\\\\ ",
        "introduction": "\\label{intro}  Extensive planet search programmes are now operating on a variety of telescopes around the world, including the 3.9m Anglo-Australian Telescope (Tinney et al. 2001,Butler et al. 2001), the 10m Keck I (Vogt et al. 2000), the Lick 3m (Fischer et al. 2001), the Leonard Euler 1.2m (Naef 2001) and ESO Coud\\'e Auxiliary Telescopes (K\\\"urster 2000) at La Silla, and the Observatoire de Haute-Provence 1.9m (Queloz et al. 1998). These programmes reach velocity precisions ranging from 15 to 3\\,\\ms. At these levels {\\em apparent} velocity variability (known as jitter), which is intrinsic to the target star itself (rather than the presence of an unseen companion), can be produced by the effects of stellar activity and rapid rotation combined with surface inhomogeneity  (Saar et al. 1998, Saar \\& Fischer 2000, Santos et al. 2000).\\\\  \\begin{figure*} \\centerline{\\psfig{file=figure1a.ps,width=2.3in}\\qquad\\qquad\\psfig{file=figure1b.ps,width=2.3in} } \\caption{Example Ca\\,\\II\\ HK spectra from UCLES. The left panel shows an observation of HD1835, indicating the locations of the Ca\\,\\II\\ H\\&K lines, and the continuum V and R regions, described in Section \\ref{sindex}. The fluctuations seen in this spectrum are due to weak absorption lines, not noise. The right panel shows an expanded region of the Ca\\,\\II\\ H line for an activity sequence from HD1835, HD16673, HD17051 to HD23429 (from top to bottom -- see Table \\ref{mytable}). Also shown are the Mount Wilson triangular Ca\\,\\II\\ bandpass we have used{\\em (solid line)}, and the Henry et al. (1996) broad Ca\\,\\II\\ bandpass {\\em (dotted line)}. All spectra have been normalised to F$_\\nu$=1.0 in the R continuum region. The spectra have been offset by 0.4,0.1,0.2,0.0 (from top to bottom) for clarity.} \\label{spectra} \\end{figure*}  The level of stellar activity in solar-type stars is commonly examined by measuring the strength  of chromospheric emission features in the cores of the Ca\\,\\II\\ H and K absorption lines. This emission is usually parametrised by the Mount Wilson HK Project's \\rhk\\ index (see Baliunas et al. (1998) and references therein). This index has been shown to be a useful indicator of the levels of jitter in F-M type stars (Saar, Butler \\& Marcy 1998). Estimates of \\rhk\\ are therefore essential to planet search programmes, since they indicate the level of velocity variability in a given star which could be erroneously ascribed to the presence of an unseen planet, rather than jitter. \\\\  Unfortunately, while extensive \\rhk\\ data are available in the Northern Hemisphere from the Mount Wilson HK Project programme (Baliunas et al. 1998) and the Vienna-KPNO Ca\\,\\II\\ H\\&K Survey (Strassmeier et al. 2000), the only large \\rhk\\ data set available in the South is that of Henry et al. (1996). Henry et al. measured chromospheric Ca\\,\\II\\ emission in more than 800 southern stars (mostly G dwarfs). However, a large fraction of the stars targeted by the Anglo-Australian Planet Search (Tinney et al. 2001, Butler et al. 2001) were not observed by Henry et al. We have therefore begun a programme of ultraviolet (UV) spectroscopy of these stars on the AAT to fill this gap.\\\\ ",
        "conclusions": "\\label{conclusion}  Ca\\,\\II\\ HK activity measurements are provided for 59 stars from the Anglo-Australian Planet Search target list -- thirty five of which are for stars not previously published. Four-metre telescope UV echelle spectroscopy is an efficient way in which to obtain high quality data for this purpose, and because of the high resolution of such data, calibration onto the standard Mount Wilson \\rhk\\ system can be carried out with greatly reduced systematic uncertainties.  \\subsection*"
    },
    "0201/astro-ph0201177_arXiv.txt": {
        "abstract": "High resolution far-ultraviolet spectra of the O-type subdwarf \\bd28\\ were obtained with the Far Ultraviolet Spectroscopic Explorer to measure the interstellar deuterium, nitrogen, and oxygen abundances in this direction. The interstellar \\di\\ transitions are analyzed down to \\lyi\\ at 920.7 \\AA. The star was observed several times at different target offsets in the direction of spectral dispersion. The aligned and coadded spectra  have high signal-to-noise ratios (S/N $=50-100$).  \\di, \\ni, and \\oi\\  transitions were analyzed with curve-of-growth and profile fitting techniques.  A model of interstellar molecular hydrogen on the line of sight was derived from \\h2\\ lines in the \\fuse\\ spectra and used to help analyze some features where blending with \\h2\\ was significant.  The \\hi\\ column density was determined from high resolution \\hst /STIS spectra of \\lya\\ to be log \\nhi\\ $=19.846 \\pm 0.035 ~(2\\sigma)$, which is higher than is typical for sight lines in the local ISM studied for D/H.  We found that D/H $=(1.39 \\pm 0.21) \\EE{-5} ~(2\\sigma)$ and O/H $=(2.37 \\pm 0.55) \\EE{-4} ~(2\\sigma)$.  O/H toward \\bd28\\ appears to be significantly below the mean O/H ratio for the ISM and the Local Bubble. ",
        "introduction": "Both curve of growth (COG) and  profile fitting techniques were used to determine \\ndi\\ and \\noi\\ in order to better understand potential sources of systematic error. Profile fitting alone was used for \\nni\\ because there are only three unsaturated and unblended  \\ni\\ lines available (the \\wl952 multiplet). The \\fuse\\ spectra were also analyzed to develop a model for interstellar \\nh2\\ for rotational levels $J=0-4$.  The comparison of measurements from multiple channels (SiC1 and SiC2) and observations (LWRS and MDRS) is a powerful tool that  improves our statistics and  helps to identify poor data quality due to instrumental effects.  In the COG method the equivalent width, \\eqw, of several lines of the same species having  different oscillator strengths are measured and their distribution in the $\\log(W_{\\lambda}/\\lambda)-\\log(Nf\\lambda)$ plane is compared with a theoretical curve of growth for a single absorption component subject to a Maxwellian velocity distribution.  The free parameters in the fit are the column density, $N$ and Doppler spread parameter, $b$. However, if there is a complex velocity distribution, $b$ is regarded as an ``effective'' Doppler spread and as such has no physical meaning. Continuum placement and equivalent width measurements are treated in independent steps and assume some a posteriori conservative error bars to take into account the uncertainties in continuum placement. The COG method is potentially vulnerable to error because of the assumption of a single-component profile. In particular, at the \\fuse\\ spectral resolution narrow, unresolved components could be concealed by broader features, which could result in underestimated column densities (Nachman \\& Hobbs 1973, Savage \\& Sembach 1991).  Jenkins (1986) showed that column densities accurate to $\\sim$10\\% can be obtained with the COG method provided that the lines are not heavily saturated $(\\tau_0 \\le 2)$ and have smooth distributions of $b$ and $\\tau$ in their components. In the case where all transitions are on the linear part of the COG, the column density is insensitive to the line-of-sight velocity structure.  One of the advantages of the COG method is that it is not too sensitive to the LSF. On the other hand, transitions blended with lines from other species should be avoided in this analysis technique. In this case, profile fitting may be more appropriate.  The profile fitting method makes the same basic assumptions as the COG: an interstellar absorption line is formed in a single component and is described by two parameters: $N$ and $b$. In addition the wavelength centroid of the line, or radial velocity, $v_r$, is also a free parameter.  The fitting technique also assumes  that each interstellar component produces absorption lines described by Voigt profiles.  We  note that in our analysis the Doppler parameter, $b$, has no unique physical interpretation.  Normally $b^2 = 2kT/m + v_{turb}^2$.  However, from the \\fuse\\ spectra alone, with spectral resolution of $\\sim 15 \\kms$, temperature ($T$) and turbulent motion ($v_{turb}$) effects on the line widths are not easily separated.  We can, however, use $b$ to set an upper limit on $T$.  Lines of sight with multiple components are often more reliably analyzed with a fitting technique than it is with the COG. In practice, to fit a multiple-component COG requires a prior knowledge of $v_r$ and $b$ for each component as well as their expected relative ratios. This information is not necessary with profile fitting because in principle it lies in the profiles.  For the same reason, the blending of adjacent lines is also easier to model with a fitting procedure.  A second significant advantage of our profile fitting analysis is that absorption lines and continua are fit simultaneously so that continuum shape and placement are treated as free parameters. We have used the profile fitting program {\\tt Owens.f} (Lemoine et al. 2001, H\\'ebrard et al. 2001). {\\tt Owens.f} is a powerful profile fitting program that simultaneously models multiple spectral lines, species, and absorption components in an arbitrary number of spectral regions, or windows.  Many quantities characterizing the fit may be treated as free or fixed parameters ($N, T, v_{turb}$, LSF, continuum placement and shape, background level, velocity offset between spectral windows, etc.). The program uses an optimized \\chisq-minimization algorithm to obtain fast convergence.  The flexibility of {\\tt Owens.f} allowed us to analyze many different cases.  As with any convergence algorithm, parameter space was explored to verify that a converged solution was the true minimum and not a local minimum far from the true solution. This was done by changing the initial parameters and/or conditions and then running the program again to reach a new converged solution. A potential disadvantage of profile fitting is that knowledge of the instrument LSF is required, and errors in the assumed LSF have the potential to bias the results.   The instrumental LSF is a key parameter in the profile fitting analysis of interstellar absorption lines.  The characterization of the \\fuse\\ LSF was incomplete at the time of this analysis (see Kruk et al. 2001).  Although there is evidence that the \\fuse\\ LSF has wings that effect the zero level of heavily saturated lines (\\S\\ref{obs}), two independent analyses (H\\'ebrard et al. 2001, Wood et al. 2001) show that modelling of weak unsaturated lines with a single Gaussian LSF and a double Gaussian LSF give very consistent results within the uncertainties of the \\fuse\\ LSF parameters (see the companion papers referenced above for further details).  For our profile fitting analyses we used a single gaussian LSF for each component modelled.   In the next subsection the interstellar \\h2\\ model is derived. The following subsections describe the COG analysis of \\di\\ and \\oi\\ and the profile fitting analysis of \\di, \\oi, and \\ni.  \\subsection{Interstellar \\h2\\ Column Densities}\\label{h2model}  The \\fuse\\ spectra of \\bd28\\ contain a large number of absorption lines from interstellar molecular hydrogen. Absorption arising from \\h2\\ rotational levels $J=0-4$ can be clearly seen along this line of sight.  Many \\h2\\ lines are blended with stellar and other ISM features. For transitions corresponding to the rotational level $J=4$ there are not enough unblended lines to derive a reliable column density using either COG or profile fitting techniques.  However, a $3\\sigma$ upper limit for \\nh2\\ $J=4$ is derived. A few \\h2\\ transitions from $J=5$ and 6 may also be present. The best candidate is Werner $0-0$ Q(5) at 1017.831 \\AA. However, the \\fuse\\ spectra of \\bd28\\ contain many unidentified stellar features and it is possible that this and other weak transitions could be stellar in origin, or if actual \\h2\\ ISM lines, blended with stellar lines. Due to these uncertainties, the column densities for $J \\ge 5$ could not be estimated. The \\h2\\ column densities for  rotational levels $J=0-3$ were measured using both profile fitting and COG methods, using lines that appeared free of blending, as judged by the shape of the absorption profile and with the help of the stellar model.   Profile fitting used only unsaturated lines, i.e. lines for which the absorption profile deconvolved with the LSF did not have a residual intensity in the line core $< 0.1$. The column density for each rotational level was determined independently.  The synthetic stellar spectrum gave guidance when selecting the \\h2\\ lines to be analyzed in order to avoid blending with stellar features.  Each spectral region (several Angstroms) was first normalized with a spline function.  A subset of this  region, a spectral window of $\\sim1$\\AA, was used by {\\tt Owens.f} where a polynomial (up to $4^{th}$ order) was fit to the continuum away from the line.  We restricted the size of the spectral windows used in the {\\tt Owens.f} analysis because of concern about small nonlinearities in the \\fuse\\ wavelength calibration on larger scales. The number of lines analyzed for each $J$ level are listed in Table \\ref{nh2table}. All the \\h2\\ lines analyzed lie away from the \\hi\\ lines. The background was set to zero for each window. The width of the Gaussian LSF was free to vary from window to window and between successive fits.  Once the best fit was found, corresponding to a minimum \\chisq, the error in the column density was calculated by fixing \\nh2\\ to a different value and computing a new optimized fit, which always had a larger \\chisq.  The difference of the new reduced \\chisq\\ from \\chisqmin\\ determined the 1, 2, 3, 4, 5, 6, and $7\\sigma$ deviations from the best fit.  The adopted $1\\sigma$ error was the average of these values to allow for any asymmetry in the $\\Delta$\\chisq\\ curve. Only an upper limit was derived for  \\nh2\\ $J=4$, the result of its weak transitions being indistinguishable from weak stellar features. \\nh2\\ and $2\\sigma$ errors derived this way are given in Table~\\ref{nh2table}.   The \\h2\\ equivalent widths were measured by directly integrating the area of several lines for each rotational level for $J=0-3$.  The continuum near each line was defined by a Legendre polynomial up to $5^{th}$ order.  \\eqw\\ of each transition is the weighted average of \\eqw\\ measured in each channel where the line was present. A  single-component Gaussian COG was fit to the data for each rotational level, and the residuals about the best-fit curve were minimized (Figure \\ref{h2cog}).   The resulting column densities are completely consistent with the profile fitting results, although with somewhat larger uncertainties.  Assuming the population of the lower rotational states of \\h2\\ is determined by collisional excitation and a Boltzmann distribution, the excitation temperatures are $T_{1,0} = 313$ K , $T_{2,1} = 294$ K and  $T_{3,2} = 296$ K. However, the similarity of $T_{i,j}$ for $J=0-3$ indicates that one or more of these assumptions may not be correct. The uncertainty in \\nh2\\ $J=0$ may permit a more typical  value of $T_{1,0} \\sim 100$ K.  The total \\h2\\ column density is log(\\nh2)$=15.13^{+0.20}_{-0.08} \\cmsq$.    \\subsection{\\di\\ and \\oi\\ Curve-of-Growth Analyses}\\label{dicog}  The high S/N of the \\fuse\\ observations of \\bd28\\ permits analysis of the \\di\\ transitions down to 916 \\AA\\ where the \\hi\\ profiles start to overlap.  However, some of the higher \\di\\ Lyman series lines (916--918 \\AA)  have strengths that vary significantly from line to line and channel to channel. These inconsistencies probably originate in the fixed pattern noise of the detector and/or in blends with weak stellar features or interstellar \\h2. Therefore, the COG analysis of \\di\\ was limited to the six transitions listed in Table \\ref{doeqw}.   Most of the \\di\\ lines (\\lyg\\ -- \\lyeta) lie in the blue wing of the corresponding \\hi\\ lines. Four out of the six transitions  fall on the linear part of the curve of growth. Hence, the velocity structure on the line of sight does not influence the final column density estimate (see \\S \\ref{diprof}).    \\subsubsection{\\di\\ Continuum Normalization}  The stellar \\hi\\ and \\heii\\ and interstellar \\hi\\ absorption profiles dominate the shape of the continuum near the \\di\\ feature.  A continuum model was constructed, where the key parameters are \\nhi, \\nh2, and $b$ for the ISM model computed with {\\tt Owens.f} and \\teff, $\\log g$ and He/H for the stellar model. We  adjusted the stellar plus interstellar model to fit the data (a small zero-point shift in the model's wavelength scale and scaling the absolute flux of the model to match the observed flux level -- see Fig. \\ref{model923}). These are discussed further below.  The stellar model has been discussed in detail in \\S \\ref{bd-model}.  The uncertainties in \\teff\\ and $\\log g$ considered in the \\nhi\\ analysis ($\\sim$5\\%) have a negligible effect on the shape of the \\hi\\ + \\heii\\ profile near the \\di\\ lines we are studying. While varying these parameters would change the goodness of the fit far from the core of the \\di\\ lines, the continuum placement  would be similar.  As to the He/H ratio, this  affects only the two strongest lines (\\di\\ \\lyg\\ and \\lyd) where \\heii\\ is readily apparent in the model.  The overall shape of the model agrees well with the spectrum, although we noted some localized discrepancies. The different spectra showed flux scale differences from the model of up to +5\\% for LWRS and +10\\% for MDRS. The latter suffered from small variations in slit transmission due to mirror motion (see Sahnow et al. 2000) and a small part of the flux was lost in this process. The MDRS flux calibration is also  less well known than that of the LWRS channel.  We modelled the interstellar \\hi\\ profiles by using {\\tt Owens.f} to fit  the relevant \\hi\\ Lyman series transitions (919 -- 972 \\AA), forcing the column density to log(\\nhi) = 19.85 and the adopted the \\nh2\\ parameters from Table \\ref{nh2table}. A one-component fit agreed well with the data. A three-component fit was also performed (cf. \\S\\ref{diprof}) but there was no improvement in the fit because in \\hi\\ these components are buried in the saturated line cores.  The purpose of this \\hi\\ model was to find a good empirical match with the data  to assist in the measurements of \\di\\ \\eqw, not to determine \\nhi.  \\subsubsection{\\eqw\\ Measurements}  The general method for measuring \\eqw\\ for the \\di\\ lines is illustrated in Figure \\ref{cogmethod}. First, the data were divided by the stellar model appropriately registered in \\wl\\ and $F_{\\lambda}$ (Fig. \\ref{cogmethod}a and d). This process sometimes produced what was judged to be a poor fit in the red wing of interstellar \\hi, especially for \\di\\,\\wl925. We believe this is  due to differences in the number, location,and strength of weak stellar lines  between the blue and red wings.  In spite of the concern about such line blanketing differences, \\eqw\\ for \\di\\, \\wl925 is consistent with the other \\di\\ transitions analyzed.  Second, the normalized spectrum was then corrected for interstellar \\hi\\ (Fig. \\ref{cogmethod}b and e) by using the \\hi\\ model profiles described above. \\eqw\\ of each \\di\\ feature was measured by fitting   a Gaussian to the residual \\di\\ absorption line (Fig. \\ref{cogmethod}c and f) using all of the profile except the far red wing of where the near-zero flux near the line core of \\hi\\ produces large residuals. Since the principal source of error in this analysis is the stellar continuum placement, we adopted a  conservative approach, scaling the continua by $\\pm$3\\% and remeasuring \\eqw.  This produced noticeably poor fits that represent extreme cases.  We assumed that these fits were $2\\sigma$ from the best fit.  Therefore, the \\ndi\\ errors include the systematic error due to continuum placement.  Table \\ref{doeqw} gives the \\di\\ \\eqw\\ values for each channel and aperture. There are four independent measurements of W for each DI transition (SiC1 and SiC2, LWRS and MDRS).  These measurements were consistent to within $1\\sigma$, except for the three cases where \\eqw\\ was omitted.  These measurements were discarded because they were inconsistent (discrepant by $\\ge$50\\%) with the other measurements of the same transition.  We concluded that these cases were probably caused by detector artifacts (fixed-pattern noise) after careful comparison of line profiles in the four aperture and channel combinations. The far right hand column in Table \\ref{doeqw} is the mean \\eqw\\ used in the COG.  The COG fit for \\di\\  is shown in Figure \\ref{diallcog}. We found log \\ndi$= 14.99 \\pm 0.10 ~(2\\sigma)$ and $b=5.2 \\kms$.  \\subsubsection{\\oi\\ Curve of Growth}  A curve of growth was constructed for \\oi, for each segment (SiC1 and SiC2, MDRS and LWRS), using the lines and $f$-values of Table \\ref{doeqw}. \\oi\\ \\wll 924.52, 936.63, and 950.88 were not used because of blending with stellar or \\h2\\ features. The general procedure to construct the curve of growth is similar to what was described above for the \\h2\\ COG analysis (\\S \\ref{h2model}). For the \\oi\\ COG we chose lines that are free of blending with other ISM or stellar features.  Unlike \\di\\ and \\h2,  \\oi\\ $f$-values may have relatively higher uncertainties in the \\fuse\\ bandpass, all of them coming from theoretical calculations (Morton 1991; Morton 2000, private communication; Verner, Barthel, \\& Tytler 1994). Unlike the pathological case of \\oi\\,\\wl1026.5 pointed out by Jenkins et al. (2000), the other \\oi\\ $f$-values appear to be self-consistent. \\eqw\\ and associated errors measured for the four spectra are listed in the Table \\ref{doeqw}.  The COG fit to the mean \\eqw\\ is shown in Figure \\ref{oiallcog} and the derived column densities are given in Table \\ref{oitable}.  A COG fit to the weighted average of \\eqw\\ for the different channels and apertures yields log \\noi$_{COG} = 16.23 \\pm 0.08 \\cmsq~(2\\sigma)$ and $b=4.5 \\kms$.  \\subsection{Profile Fitting Analysis}\\label{diprof}  A profile fitting analysis of \\di, \\ni, and \\oi\\ column densities was performed to provide a comparison with the results of the COG analysis in the previous subsection. The general technique is identical to that discussed in \\S \\ref{h2model}.  Since {\\tt Owens.f} was used to fit \\di, \\ni, and \\oi\\ simultaneously, the following discussion also applies to \\noi\\ and \\nni.   The five \\di\\ transitions analyzed  (\\lye, \\lyz, \\lyeta, \\lyth, and \\lyi)  are not saturated. With the use of a fitting procedure, the choice of the lines to analyze is different than with the COG technique. An example is \\lyz\\ which was not used in the COG analysis because it is  blended with \\h2\\ $J=1$ and $J=2$. By adopting the \\h2\\ solution (Table \\ref{nh2table}) as a constraint in the fit, this transition is now part of the fitting analysis and easily deblended from the \\h2.  We note that the errors in \\nh2\\ are not significant for the \\ndi\\ analysis because the \\h2\\ lines blended with \\lyz\\ are on the flat part of the COG. The two strongest \\di\\ transitions used for the COG (\\lyg\\ and \\lyd) were omitted from the profile fits because they are mildly saturated ($\\tau_0 \\sim 2$).  Uncertainties in our knowledge of the \\fuse\\ LSF could influence the results derived from profile fits to such lines. In any case, \\ndi\\ is well constrained by the other five transitions.  Although the continuum close to the \\oi\\ lines is very smooth and easily fit with a low-order polynomial, some of the  \\oi\\ transitions used in the COG are indeed  saturated and for this reason they were not included in the \\noi\\ profile fitting analysis using {\\tt Owens.f}. Only three \\oi\\ transitions were used to determine \\noi: 919.92 \\AA, 925.45 \\AA, and 930.26 \\AA. The strongest of these has $\\tau_0 \\sim 1$.  There are many \\ni\\ transitions in the \\fuse\\ range and four multiplets are  detected : \\wll952, \\wll953, \\wll964, and \\wll1134. The 952 \\AA\\ triplet is free of any blending with \\h2\\ or ISM atomic lines and seems clear of stellar features. Since all the lines of this triplet are still on the linear part of the curve of growth, deriving the column density is straightforward. This is not the case for the \\wll953, \\wll964, or  \\wll1134. These multiplets ets are all  saturated and much less sensitive to $N$ than they are to $b$. The 964 \\AA\\ triplet is the next strongest after \\wll952, but these lines are blended with several \\h2\\ lines, \\pii\\, \\wl963.8, and with unidentified stellar features.  With only one multiplet (\\wll952) we did not perform a COG analysis of \\ni.  The \\ni\\ $f$-values used were provided by D. Morton (2000, private communication).  The {\\tt Owens.f} analysis of the \\di, \\ni, and \\oi\\ lines allow many quantities that characterize the fits to be treated as free parameters.  These include the wavelength and flux zero points of each profile, the polynomial coefficients for each spectral window, the LSF parameters (discussed below), the column densities $N$, and Doppler parameters $b$.  In total, there were about 40 spectral windows and about 1500 degrees of freedom to fit 11 transitions.  The \\h2\\ column densities (Table \\ref{nh2table}) were adopted as a fixed constraint. We fit the five lines of \\di, three of \\oi, and three of \\ni\\ in the four data sets simultaneously, providing a global solution.   The continuum for the \\di\\ line profiles was determined by a different approach than was used for the \\di\\ \\eqw\\ measurements. There is sufficient information immediately adjacent to the \\di\\  features to establish the continuum with a polynomial without normalizing the spectrum or modelling the stellar and interstellar \\hi. The shoulder of the blue wing of the interstellar \\hi\\  profile, excluding the \\di\\ feature, was fit  by a $4^{th}$ order polynomial.   Examples of the continuum and \\di\\ profile fits are shown in Figure \\ref{difits}. The effective $b$ is not tightly constrained in this analysis because the true shapes of the \\di\\ profiles are unresolved for these weak lines.  The effect of simultaneously modelling the interstellar \\hi\\ and \\di\\ profiles and fitting a polynomial to the stellar contribution alone was also examined. The line-of-sight velocity structure  becomes critical when higher Lyman series of \\di\\ and \\hi\\ are fit together. In this case for \\bd28\\ the effective $b$ is dominated by \\hi\\ and is  larger than the \\di\\ effective $b$, much larger than the expected $\\sqrt 2$ difference between H and D.  This is probably the result of very weak \\hi\\ components not visible in \\di\\ or any other species that become a factor in the width high Lyman series profiles of \\hi. Overestimating $b$ for \\di\\  leads to an anomalously low \\ndi.   When \\di\\ and \\hi\\ are fit together a larger fitting window ($\\Delta\\lambda\\sim$4 \\AA) is also required, forcing \\di\\ to be shifted from its true center as a result of small residual nonlinearities in the \\fuse\\ wavelength calibration.  Due to the introduction of these systematic errors we rejected this fitting approach in favor of the techniques described in the previous paragraph and illustrated in Figure \\ref{difits}.  The influence of the LSF assumptions and the velocity structure were carefully examined. Three cases were considered in the profile fitting analysis.   One way of investigating the sensitivity of the fit to the assumed LSF was to let {\\it Owens.f} adjust it as part of the \\chisq\\ minimization. For Case 1 we considered a fixed LSF associated with each channel/aperture combination, ignoring the possible variations with  wavelength.  The choice of the LSFs was made after several {\\tt Owens.f} test runs with different parameters. In Case 2, we assumed a LSF modeled by a single Gaussian with FWHM that varied freely with wavelength and from channel to channel. Case 3 examined the effect of line-of-sight velocity structure as derived from STIS \\feii\\ profiles on the derived column densities.  The continuum was determined in all three cases by the technique illustrated in Figure \\ref{difits}.  For all three cases, the LSF was modelled as a single Gaussian for reasons discussed at the end of \\S\\ref{dnointro}.  A free LSF (Case 2) clearly improved the fit, as indicated  in Table \\ref{dnotable}, compared with the fixed LSF (Case 1). Only a few spectral windows showed  LSF width variations as large as 20\\% ($\\le2$ pixels).  Such variations may be the result of fixed pattern noise that is only partly corrected by means of the FP-split procedure, or from weak, unresolved photospheric features.  For Case 2  the reduced \\chisq\\ was slightly lower than before but the \\ndi\\ and \\noi\\ were unchanged.  The fit of the \\ni\\ lines, fit simultaneously with \\di\\ and \\oi, showed slight differences but are within the $2\\sigma$ uncertainty. Although  the LSF calibration as a function of \\wl, channel, and aperture is not yet well determined (Kruk et al. 2001), we accounted for this by considering two extreme assumptions (free vs. fixed LSF) and adopting an average column density that includes this systematic effect in the errors.  For Case 3 we used high resolution STIS E230H spectra to provide resolved \\feii\\, \\wll2344, 2382 line profiles, the only \\feii\\ lines available in the existing archival high resolution spectra of \\bd28. Profile fits to the STIS spectra alone found three components at $v_{hel}= -2.4, -9.7$, and $-17.0 \\kms$ (listed in order of decreasing \\feii\\ column density) and a total \\feii\\ column density of $\\log N$(\\feii)$= 13.81 \\cmsq$ (see Figure \\ref{feiifit}).  We also performed a simultaneous fit of \\fuse\\ and STIS data using {\\tt Owens.f}. This fit found the velocity separations of the three components to be 3.5 and 10.9 $\\kms$ and $\\log N$(\\feii)$=13.79 \\cmsq$, compared with 7.3 and 7.3 $\\kms$ for the STIS data alone. These differences $(3-4 \\kms)$ reflect the level of uncertainty of these fits.  For the Case 3 fits, the relative ratios of each component for the neutral species was very different from that in  \\feii. Specifically,  the weakest (bluest) \\feii\\ component had no contribution from \\di, \\ni, or \\oi. This could be due to ionization effects or iron might be depleted in different ways in different environments, while \\ni\\ and \\oi\\ are only lightly depleted (Savage \\& Sembach 1996). It appears  that the fit of the neutral species suffered from the use of the \\feii\\ lines, showing a higher reduced \\chisq. We believe that the neutral species are not  fit well by the \\feii\\ velocity structure because the $b$-values and/or the velocity structure appears to be different. Since we do not resolve the component structure with \\fuse, we rely on the one component solution for \\di, \\ni, and \\oi.   The column density errors were calculated by computing new fits for different fixed values of $N$ away from the best fit column density and deriving  $\\Delta$\\chisq\\ as a function of $N$.  The $3\\sigma$ error is given by $\\Delta$\\chisq = 9.  The average gives log \\ndi\\ $= 14.99 \\pm 0.03 \\cmsq (2\\sigma)$.  The COG and profile fitting analyses give the same \\di\\ column density and have error estimates that differ only slightly at the $2\\sigma$ level. It is reassuring that both approaches give results within 5\\% of each other.  The errors  from the COG analysis are more conservative than the ones coming from the profile fits. The reason for this is that moving the continuum by $\\pm3\\%$, as was done for the COG, we have probably overestimated the errors. Yet they seem to fairly reflect  the scatter  seen in the individual \\eqw\\ measurements.   Combining the results,  we adopted a straight average of \\ndi\\ from the COG and profile fitting analyses and arrive at the final \\di\\ column density log \\ndi\\ $= 14.99 \\pm 0.05 \\cmsq ~(2\\sigma)$ for the \\bd28\\ line of sight.    Figure \\ref{oifits} compares the \\oi\\ fits for SiC2 spectra from the LWRS and MDRS observations. Combining the results from the different LSF approaches (see Table \\ref{dnotable}) gives log \\noi$=16.21\\pm 0.10 \\cmsq~(2\\sigma)$ from the profile fitting analysis.   Most of the \\oi\\ lines used in the profile fitting analysis are close to saturation which may make both this and the COG method sensitive to the velocity structure. We investigated this possibility by analyzing the effect of a three-component velocity structure on the COG analysis. The goal of this investigation was to see if a one-component COG  could still be applied with success on this line of sight. We computed a model for the line of sight for \\di\\ and \\oi\\ lines with three components as seen in the STIS \\feii\\ profiles. Performing a one-component COG analysis of the model line profiles yielded \\di\\ and \\oi\\ column densities within 5\\% of the model and an average effective $b$ of 5.8 $\\kms$. Although the velocity structure found in \\feii\\ is not identical to that for the neutral species, it is the best available. For the \\bd28\\ line of sight   a one-component COG gives reliable results. We find that the effective $b$ from the analysis of the three-component simulated profiles is consistent with the effective $b$ from the analysis of the actual data. The final value adopted for the \\oi\\ column density is  log \\noi $= 16.22\\pm 0.10 \\cmsq~(2\\sigma)$.  Examples of the fits to the \\ni\\,\\wll952 lines are shown in Figure \\ref{nifits} and the results listed in Table \\ref{dnotable}. The mean fit gives log \\nni $=15.55 \\pm 0.13 ~(2 \\sigma)$. ",
        "conclusions": "The abundances derived in the previous sections are summarized in Table \\ref{dhtable} along with their ratios.  The $2\\sigma$ errors in the abundance ratios are those from the individual column densities combined in quadrature. We find that D/H $=(1.39\\pm 0.21) \\EE{-5}~(2\\sigma)$ toward \\bd28.  G\\\"olz et al. (1998) observed \\bd28\\ with the T\\\"ubingen echelle spectrograph ($\\lambda/\\Delta\\lambda \\sim$ 10,000) on the \\orf\\ mission and found D/H $=(0.8^{+0.7}_{-0.4}) \\EE{-5} (1\\sigma)$. Our result is at the upper end of their range.   The \\fuse\\ result for \\bd28\\ is just marginally consistent with the higher D/H value found for the more distant sight line toward Feige 110 (D/H $=(2.1\\pm 0.70) \\EE{-5}$) by Friedman et al. al. (2001).  The total neutral gas column toward \\bd28\\ is very similar to that toward \\gvel, \\zpup, and \\dori\\ (at distances of 250--500 pc), the sightlines studied with IMAPS to measure D/H (Jenkins et al. 1999, Sonneborn et al. 2000).  D/H toward \\bd28\\ is in the middle of the range of values found in the IMAPS studies.  Our result for D/H toward \\bd28\\ is similar to values found on lines of sight in the Local Bubble by \\fuse\\ (Moos et al. 2001) and \\hst\\ (Linsky et al. 1998), even though \\bd28\\ lies just outside the Local Bubble. The total \\hi\\ column density toward \\bd28\\ is  log \\nhi = 19.85, much higher than the other LISM sight lines studied thus far by \\fuse\\ for D/H (except for Feige 110). The \\nai\\ contours for the LISM (Sfeir et al. 1999) show that \\bd28\\ is located in a region of higher column density just beyond the edge of the Local Bubble, with much lower densities in the intervening $\\sim100$ pc. The \\hi\\ column toward \\bd28\\ is large enough that the local cloud is a negligible contribution.  Our abundance measurements are therefore dominated by the gas $\\sim100$ pc from the Sun.   As discussed by Moos et al. (2001), there is now considerable evidence for a constant D/H within the Local Bubble.  Beyond $\\sim100$ pc, however, or more specifically outside the Local Bubble and log \\nhi\\ $ge 19.5$, there appears to be a considerable spread in D/H as larger amounts of gas are probed.  The total \\h2\\ column density is $\\sim 1.3 \\EE{15} \\cmsq$, and the molecular fraction $f({\\rm H_2}) \\sim 3.8 \\EE{-5}$.  This low value of $f({\\rm H_2})$ is typical of lower \\hi\\ column density ISM sight lines (Tumlinson et al. 2001).  The abundance  of \\h2\\ at $\\sim100$ pc indicates that some molecular gas is present, but the low column densities and high excitation temperature are consistent with a low density environment where UV pumping is probably significant.  The profile fitting and COG analyses for \\di\\ and \\oi\\ consistently found an effective $b$-value of $\\sim 5 \\kms$.  Although \\fuse\\  lacks the spectral resolution to sepatate thermal and non-thermal effects on the neutral gas line profiles, we can set an upper limit to the temperature.  For \\di\\ $b=5.0 \\kms$ corresponds to $T\\le 6000$ K.  For \\oi, however, the same $b$ corresponds to $T\\le$40,000 K.  There must either be a significant non-thermal contribution to the line broadening of the neutral species, or the component structure on the \\bd28\\ line of sight has a significant effect of the line shapes.  Very high resolution spectra of species such as \\nai\\ are needed to determine the component structure, $T$ and $v_{turb}$ of the neutral gas as a function of velocity.  The O/H ratio toward \\bd28\\ is $(2.37\\pm 0.55) \\EE{-4} ~(2\\sigma)$, a value that is anomalously low with respect to the Local Buble and the ISM.  The \\bd28\\ O/H is marginally consistent with the mean O/H ratio  in the ISM ($<$O/H$>_{ISM}$) of $(3.19\\pm0.28) \\EE{-4} ~(2\\sigma)$ found by Meyer et al. (1998) in a  study of \\oi\\ \\wl1355 on 13 sight lines toward OB stars at  distances of 130 -- 1500 pc.  Recently, however, Meyer (2001) updated $<$O/H$>_{ISM}$ to $(3.43\\pm0.28) \\EE{-4}$ as a result on a revised $f$-value for the \\wl1355 transition.   The \\fuse\\ O/H ratio toward \\bd28\\ is well outside the estimated uncertainties in the revised $<$O/H$>_{ISM}$.  In the Local Bubble $<$O/H$>_{LB}=(3.9\\pm0.3) \\EE{-4}~(1\\sigma)$ (Moos et al. 2001),  slightly higher than Meyer's revised value for the ISM, although consistent within their respective uncertainties.  Moos et al. (2001) also finds the same value of the mean O/H when the more distant sight line to Feige 110 is included. Of the five sight lines for which O/H is measured by \\fuse, only \\bd28\\ gives a value that disagrees with $<$O/H$>_{ISM}$ and $<$O/H$>_{LB}$. If correct, the gas just outside the Local Bubble has a low O abundance.   This implies that the gas toward \\bd28\\ beyond the Local Bubble would appear to have a different chemical history than that sampled in the Meyer et al. study.  We have searched for other possible explainations for the low O/H ration toward \\bd28.    \\oi\\ is not vulnerable to ionization effects (Jenkins et al. 2000), so the \\oi-to-\\hi\\ ratio should accurately reflect the total O/H abundance ratio.  If the error is in the \\hi\\ column density, then correcting O/H to obtain agreement with $<$O/H$>_{ISM}$ would result in a large D/H ratio ($\\sim2.0 \\EE{-4}$). Such a D/H value for \\bd28\\ would be similar to that found for Feige 110 (Friedman et al. (2001) and \\gvel\\ (Sonneborn et al. 2000).  However, a $\\sim$40\\% increase in \\nhi\\ would be inconsistent with the available spectra.  The D/O ratio for \\bd28\\ appears to be high (Moos et al. 2001, H\\'ebrard et al. 2001), arguing for \\noi\\ as the anomalous quantity. Agreement with $<$O/H$>_{ISM}$ would require a $\\sim40$\\% increase in \\noi.  One or more of the \\oi\\ lines analyzed might actually be saturated, even though our analysis shows (Fig. \\ref{oifits}) that the three \\oi\\ lines are unsaturated.  Line saturation, possibly as the result of narrow (cold gas) components  would result in an underestimated column density. The effect of potential line saturation in \\noi\\ was examined by excluding the strongest \\oi\\ line (\\wl930) from the fitting analysis, leaving only the two weaker \\oi\\ lines.  For this case we found log \\noi\\ = 16.28, which is within $2\\sigma$  of the best fit.  We cannot determine whether this difference in $N$ is the result of saturation, fixed pattern noise, LSF uncertainties, component structure, or some other cause. If this were the correct \\oi\\ column density, the O/H ratio would be $2.72 \\EE{-4}$, still well below $<$O/H$>_{ISM}$ and $<$O/H$>_{LB}$.   The N/H ratio toward \\bd28\\ is $(5.08\\pm 1.66) \\EE{-5} ~(2\\sigma)$, a value that is marginally consistent with the mean ISM value of N/H $=(7.5\\pm0.8) \\EE{-5} ~(2\\sigma)$ found for six OB stars by Meyer et al. (1997).  However, the \\nii\\ column density could be significant (see below).  If so, then the N/H ratio toward \\bd28\\ could well be consistent with the mean ISM result.  In the local ISM, the abundance of \\ni\\ has been found to be reduced by ionization effects on some sight lines (Jenkins et al. 2000, Moos et al. 2001). \\ni\\ has a photoionization cross section larger than that of \\hi\\ and a large fraction of it can be ionized. There are only two \\nii\\ transitions in the \\fuse\\ bandpass (\\wl915.6 and \\wl1085.5) and both of them are heavily saturated on all but the lowest column density sight lines. We derived a lower limit of log \\nnii\\ $>14.2$ using both optical depth technique and profile fitting. Since we know that some \\nii\\ is associated with the neutral component, we estimate an upper limit assuming that the  $b$-value for \\nii\\  is greater than that of the neutral gas. These assumptions give an upper limit of log \\nnii\\ $<15.7$. We cannot rule out the possibility that an apreciable fraction of the nitrogen on the sight line toward \\bd28\\ could be singly ionized.    Objects like the hot subdwarf \\bd28\\ are important targets for studying \\di\\ and \\oi\\ in the ISM with \\fuse\\ because they can sample regions of space beyond that accessible with white dwarfs ($d < 100$ pc) and closer that that sampled by the lightly-reddened O stars observable by \\fuse\\ ($d>1$ kpc).  Most of the nearer O stars far exceed the \\fuse\\ brightness limit.  High-quality spectra of the type shown in Figures \\ref{bdspec} and \\ref{lyspec} are needed to understand these objects, and the observational techniques to obtain them with \\fuse\\ are now available. Over time issues like stellar continuum placement and stellar line identifications may be better understood, allowing even more precise interstellar abundance measurements."
    },
    "0201/astro-ph0201341_arXiv.txt": {
        "abstract": "{ Some recent observations of the abundances of s-process, r-process, and $\\alpha$ elements in metal-poor stars have led to a new scenario for their formation. According to this scenario, these stars were born in a globular cluster and accreted the s-process enriched gas expelled by cluster stars of higher-mass, thereby modifying their surface abundances. Later on, these polluted stars evaporated from the globular cluster to constitute an important fraction of the current halo population. In addition, there are now many direct observations of abundance anomalies not only in globular cluster giant stars but also in subgiant and main-sequence stars. Accretion provides again a plausible explanation for (at least some of) these peculiarities. Here we investigate further the efficiency of the accretion scenario. We find that in concentrated clusters with large escape velocities, accretion is very efficient and can indeed lead to major modifications of the stellar surface abundances. ",
        "introduction": "Strong correlations between the r-process and s-process element abundances and the $\\alpha$-element abundances in field metal-poor stars have been reported (\\cite{jehin98}, 1999), separating these stars into two sub-populations. The r-process elements correlate linearly with the $\\alpha$-elements, with a clumping at the maximum value of $[\\alpha/Fe]$. The s-process elements, on the other hand, exhibit a more complex behavior when plotted against the $\\alpha$ elements, and form a ``two-branches diagram'', which we show schematically in Fig.~1. We emphasize here that the observed stars being dwarf stars, they cannot have synthetized these s-elements in their interior.  \\begin{figure}[h] \\centerline{\\psfig{figure=ease.ps,width=7.7cm}} \\caption[] {The ``two-branches diagram'' (see \\cite{jehin99} for details).} \\label{figease} \\end{figure}  In order to explain these observations the EASE scenario has been developed (\\cite{jehin99}; \\cite{parmen99}, \\cite{parmen00a}, \\cite{parmen00b}) which links the metal-poor field halo stars to the halo globular clusters (GCs). According to this model, the field halo stars are born in GCs. The evolution of the GCs is separated into two phases, directly connected to the two subpopulations now observed in the halo. During the first phase, ``first generation'' stars are formed. The initial mass function in the very metal-poor medium favors the high masses, with $3\\; \\Msun < M <60\\; \\Msun$. The most massive of these stars quickly evolve to become supernovae, ejecting $\\alpha$- and r-process elements into the surrounding intracluster medium. The explosions trigger the formation of an outwardly expanding supershell. The primordial ISM gas is swept into the shell with the supernovae ejecta, and the resulting gas is enriched in $\\alpha$- and r-process elements. A new burst of star formation occurs in the shell, where the density is very high. This triggered star formation presumably has a Salpeter-like initial mass function. A proto-globular cluster is born. If the gravitational potential is too weak, it gets disrupted before having the chance of forming a GC, and the stars become field halo stars of the PopIIa, as defined on Fig.~\\ref{figease}. The metallicity reached in the proto-GC depends on the number of supernovae which have exploded before the proto-GC gets disrupted. The $\\alpha$-element abundances are fixed by that time as well as by the initial mass function (supernova progenitors of different masses produce different relative amounts of $\\alpha$ elements). The proto-globular cluster may in some cases survive this supernova phase. If the stars fall back under the effect of gravity, they form a GC. Detailed dynamical calculations have been performed (\\cite{parmen99}) which show that proto-globular clusters can indeed survive the supernova phase and metallicities such as those observed today in GCs can be achieved in this scenario. It has also been shown that the EASE scenario, with its two separate star formation events, does not contradict observational constraints such as the narrow metallicity range found in any given cluster.  During the second phase, intermediate-mass (second generation) stars evolve, becoming asymptotic giant branch (AGB) stars. Those stars produce s-elements, which are brought to the surface through the so-called third dredge-up occurring after the thermal pulses. The s-elements are then ejected into the surrounding intracluster medium by stellar winds. The intracluster gas is therefore enriched in s-elements, while the $\\alpha$- and r-process element abundances have been fixed at their highest value after all massive stars have exploded. If the cluster can retain this enriched gas, it can be accreted by the lower-mass stars which are still on the main sequence. If the accreted material does not get mixed within the entire star but only in its convective zone, which contains about 1\\% of the total stellar mass for stars still near the main-sequence turn-off of present GCs, then even a small amount of accretion can lead to an appreciable enrichment of the stellar atmosphere. When the star evolves towards the giant stage, its convective zone becomes much larger, and the s-element enrichment is diluted considerably. We point out here that some mechanisms could be at work to already trigger the dilution of the accreted matter below the convective zone of {\\it main-sequence} stars (\\cite{Proffitt89}; \\cite{ProffittMichaud89}), thus diminishing the impact on the surface abundances of the deep convective envelope developing in red giant stars. We come back to this issue in more detail in Sect.~2. The accreting star can be ejected out of the GC at any time, by evaporation or by tidal forces possibly leading to the total disruption of the cluster while crossing the disk. There is indeed now observational evidence that GCs lose stars: two cannon-ball stars in 47~Tuc (\\cite{meylan91}; \\cite{jorissen01}; see also \\cite{capaccioli93}; \\cite{piotto97}, and references therein), and, above all, the recent observations of ``tidal tails'' around {\\bf many} GCs (\\cite{grillmair95}; \\cite{leon00}; \\cite{odenkirchen01}). Stars enhanced in s-process elements by accretion and then stripped from GCs form the Pop IIb halo.  The hypothesis that the gas ejected by the intermediate-mass stars in GCs can be accreted by the other (lower-mass) cluster stars has been examined before. There is strong observational and theoretical evidence that stars reaching the AGB phase lose a large amount of mass. Several studies have been devoted to the fate of this gas in clusters (\\cite{scott75}; \\cite{faulkner77}; \\cite{vandenberg77}; \\cite{vandenberg78}; \\cite{scott78}; \\cite{faulkner84}; \\cite{faulknercol84}; \\cite{smith96}). Stellar ejecta in GCs with shallow potential wells can leave the cluster via a smooth wind-like outflow. Faulkner \\& Freeman (1977) and VandenBerg \\& Faulkner (1977) find steady-state time-independent flow solutions in clusters of $10^5\\,\\Msun$. However, they show that in GCs with deep enough potential wells, the gas can accumulate into the cluster, forming a central reservoir with a radius comparable to the GC core radius. Faulkner \\& Coleman (1984) show that a small number of low-velocity low-mass stars in the cluster core can in this case accrete enough matter to form $10\\;\\Msun$ black holes, which can be considered as an extreme case. All those studies looked at what happens in present-day GCs. Smith (1996) addressed the question whether such winds might have been possible within young GCs, during epochs of much higher stellar mass-loss rates. Many GCs have high enough escape velocities to retain at least some of the stellar ejecta. A substantial amount of intracluster material could have been acquired when the turn-off mass was about $5\\,\\Msun$. In less tightly-bound clusters, the stellar ejecta are lost from the cluster either stochastically or through a continuous wind. In the cases where the stellar ejecta are retained in the cluster, the gas could be accreted by other cluster stars, thereby modifying their surface composition. Several models estimating the efficiency of the accretion scenario have appeared in the literature (see e.g. \\cite{dantona83}; \\cite{smith96}). A quantitative model investigating the recycling of nova ejecta has been presented by Smith \\& Kraft (1996). Qualitative discussions of the accretion model have been presented by Bell et al. (1981), Norris \\& Da Costa (1995), and Cannon et al. (1998).  In this paper, we study the process of accretion onto GC stars from a central reservoir of gas. In particular we take into account the evolution of this process with time. Indeed the rate of mass ejection into the cluster is strongly time-dependent, as it is highest when intermediate-mass stars reach the mass-losing stage, decreasing very rapidly as the turn-off mass decreases. In Sect.~2, we examine the possible links between halo stars and GCs. We address the problem of the ``missing'' intracluster gas in Sect.~3. We present some observational signatures that may possibly be related to accretion by cluster stars in Sect.~4. In Sect.~5 we present the method used to calculate the efficiency of accretion of gas by GC stars. The results are summarized in Sect.~6. Finally Sect.~7 contains the conclusions. ",
        "conclusions": "It was  shown in this paper that accretion by low-mass GC stars of the gas ejected in the intracluster medium by moderately massive stars may be quite efficient. If enough mass is accreted, it can lead to major alteration of the stellar surface composition. This supports the EASE scenario and provides a plausible explanation for the lack of intracluster gas, and for some of the abundance anomalies observed in GC stars."
    },
    "0201/hep-ph0201262_arXiv.txt": {
        "abstract": "We examine how constraints can be placed on the neutrino component of dark matter by an accurate measurement of neutrinoless double beta ($0\\nu\\beta\\beta$) decay and the solar oscillation amplitude. We comment on the alleged evidence for $0\\nu\\beta\\beta$ decay. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0201/astro-ph0201427_arXiv.txt": {
        "abstract": "I review recent progress in numerically simulating the formation and evolution of galaxies in hierarchically clustering universes. Special emphasis is given to results based on high-resolution gas dynamical simulations using the N-body hardware integrator GRAPE. Applications address the origin of the spin of disk galaxies, the structure and kinematics of damped \\Lya systems, and the origin of galaxy morphology and of galaxy scaling laws. ",
        "introduction": "Motivated by the increasing body of evidence that most of the mass in the universe consists of invisible ``dark'' matter, and by the particle physicist's inference that this dark matter is made of exotic non-baryonic particles, a new and on the long run more fruitful approach to study the formation of galaxies has been developed: rather than to model the formation and evolution of galaxies from properties of present day galaxies, it is attempted to prescribe a set of reasonable initial conditions. The evolution of galaxies is then modeled starting from these initial conditions. Physical processes are taken into account that are considered to be relevant such as gravity, gas dynamics, radiative cooling and star formation. The outcome at different epochs is then confronted against observational data.  One scenario that has been extensively tested in that way is the model of hierarchical clustering, currently the most successful paradigm of structure formation in the universe. In this scenario, structure grows as objects of progressively larger mass merge and collapse to form newly virialized systems. The probably best know representative of this class of models is the {\\sl Cold Dark Matter} (CDM) scenario. The initial conditions consist of the cosmological parameters ($\\Omega, \\Omega_{\\rm baryon}, \\Lambda, H_0$) and of an initial fluctuation spectrum such as the CDM spectrum. The remaining free parameter, the amplitude of these initial fluctuations, is calibrated by observational data, \\eg, the measured anisotropies of the microwave background. Over the past few years, limits on the values allowed for these parameters have been consistently refined by improved observational techniques and theoretical insight, and it is widely accepted that a new ``standard'' model has emerged as the clear front-runner amongst competing models of structure formation. This $\\Lambda$CDM model envisions an eternally expanding universe with the following properties (Bahcall \\etal 1999): (i) matter makes up at present less than about a third of the critical density for closure  ($\\Omega_0 \\approx 0.3$); (ii) a non-zero cosmological constant restores the flat geometry predicted by most inflationary models of the early universe ($\\Lambda_0=1-\\Omega_0\\approx 0.7$); (iii) the present rate of universal expansion is $H_0 \\approx 70$ km s$^{-1}$ Mpc$^{-1}$ ($h=H_0/100$ km s$^{-1}$ Mpc$^{-1} \\approx 0.7$); (iv) baryons make up a very small fraction of the mass of the universe ($\\Omega_b \\approx 0.019 \\, h^{-2} \\ll \\Omega_0$); and (v) the present-day {\\sl rms} mass fluctuations on spheres of radius 8 $h^{-1}$ Mpc is of order unity ($\\sigma_8 \\approx 0.9$).  The hierarchical structure formation process in this $\\Lambda$CDM scenario is illustrated in Figure 1, which depicts the growth of structure within a $32.5/h\\,$Mpc box between redshifts nine and zero.  The $\\Lambda$CDM model is consistent with an impressive array of well-established fundamental observations such as the age of the universe as measured from the oldest stars, the extragalactic distance scale as measured by distant Cepheids, the primordial abundance of the light elements, the baryonic mass fraction of galaxy clusters, the amplitude of the Cosmic Microwave Background fluctuations measured by COBE, BOOMERANG, MAXIMA and DASI, the present-day abundance of massive galaxy clusters, the shape and amplitude of galaxy clustering patterns, the magnitude of large-scale coherent motions of galaxy systems, and the world geometry inferred from observations of distant type Ia supernovae, among  others.  \\begin{figure*} \\epsfig{file=figs/fig1.eps,height=3.2cm} \\caption[]{Time sequence of structure formation in a hierarchical clustering universe, here for the so-called $\\Lambda$CDM model. The four snapshots correspond (from left to right) to redshifts of 9, 3.5, 1 and 0, respectively. The side length of the simulation box is $32.5\\,h^{-1}$Mpc comoving.} \\end{figure*}  The hierarchical build-up is also thought to determine the morphology of a galaxy, most noticeably the difference between disk like systems such as spiral galaxies (some of them barred) and spheroidal systems such as elliptical galaxies and bulges. This picture envisions that whenever gas is accreted in a smooth fashion, it settles in rotationally supported disk-like structures in which gas is slowly transformed into stars. Mergers, however, convert disks into spheroids. The Hubble type of a galaxy is thus determined by a continuing sequence of destruction of disks by mergers, accompanied by the formation of spheroidal systems, followed by the reassembly of disks due to smooth accretion. This picture of a hierarchical origin of galaxy morphology has been schematically incorporated in so-called semi-analytical galaxy formation models used to study the evolution of the galaxy population, but its validity in a cosmological setting has just recently been directly demonstrated (Steinmetz \\& Navarro 2002).  Numerical simulations have been an integral part in the detailed analysis of the virtues of the CDM scenario. Only numerical techniques can account for the highly irregular structure formation process and for at least some of the complicated interaction between gravity and other relevant physical processes such as gas dynamical shocks, star formation and feedback processes. Simulations also provide the required interface to compare simulations with observational data and are able to link together different epochs. While simulations of structure formation on the larger scales have mainly used large massively parallel supercomputers, studies how individual structures such as galaxies or clusters of galaxies form in the $\\Lambda$CDM scenario have heavily used special purpose hardware like the GRAPE (=GRAvity PipE) family of hardware N-body integrators (Sugimoto \\etal 1990).  In this review I will concentrate on some examples, how high resolution gas dynamical simulation using special purpose hardware have illustrated some successes but also also some failures of $\\Lambda$CDM scenario in reproducing structures as observed on the scales of galaxies.  \\begin{figure} \\mbox{\\hskip0.7cm\\epsfig{file=figs/fig2.eps,width=12cm}} \\caption{\\label{damped}Right: Color map of the column density distribution in a 60\\,kpc box around a DLAS. Black correspond to \\HI\\ densities $\\log n (\\HI) > 1.5$), light gray to $\\log n (\\HI) \\approx -3$). Arrows indicate the velocity field. The solid line corresponds to the line-of-sight (LOS). The lower left plot shows, the velocity field along the LOS, the upper left plot the absorption line  in C{\\scriptsize IV} 1548 (top) and Si{\\scriptsize II} 1808 (bottom). For readability, C{\\scriptsize IV} has been displaced  by 0.5 in flux.} \\end{figure} ",
        "conclusions": "I presented some results of recent efforts to model the formation and evolution of galaxies in a hierarchical structure universe using high-resolution computer simulations. I demonstrated that only numerical simulation can take full account of the dynamics of the formation process and the complicated interplay between different physical processes such as, \\eg, accretion and merging, star formation and feedback, photo heating and radiative cooling. Observational data can easily be misinterpreted if these effects are not properly included.  For example, the apparent inconsistency of hierarchical structure formation models and the kinematics of high-$z$ damped \\Lya absorption systems could be easily solved by properly accounting for the complicated non-equilibrium dynamics of galaxies in the process of formation. Interestingly, the opposite is the case for the example of the galaxy spin. The simplifying assumption of collapse under conservation of angular momentum combined with success of this model in explaining the sizes of disk galaxies may lure someone to consider the origin of the sizes of disk galaxies as being solved. However, at closer inspection using numerical simulation, the physics behind the sizes of disk galaxies may be far more complicated. Maintaining the hierarchical build-up of galaxies and simultaneously avoiding substantial exchange of angular momentum from the gas to the dark matter due to mergers appears to be a major challenge to the scenario, and it cannot be excluded that the angular momentum crisis may finally lead to substantial revisions of the hierarchical structure formation scenario."
    },
    "0201/astro-ph0201226_arXiv.txt": {
        "abstract": "We study the structural characteristic  of the variable DA white dwarf G117B-15A  by applying the  methods of  asteroseismology.  For  such a purpose, we  construct white  dwarf evolutionary models  considering a detailed and  up - to -  date physical description as  well as several processes  responsible for  the occurrence  of element  diffusion.  We have considered  several thickness  for the outermost  hydrogen layer, whereas  for the  inner  helium-, carbon-  and  oxygen-rich layers  we considered realistic  profiles predicted by calculations  of the white dwarf progenitor evolution.  The stellar masses we have analysed cover the mass range of $0.50 \\leq {\\rm M}_{*}/M_{\\odot} \\leq 0.60$.  The evolution of each of  the considered model sequences were followed down to very low effective  temperature; in particular, from 12500K on we computed  the dipolar,  linear, adiabatic oscillations  with radial order $k= 1, \\cdots, 4$.  We find that asteroseismological results are not univocal regarding mode  identification for the case of G117B-15A. However,   our  asteroseismological   results   are  compatible   with spectroscopical data only if the  observed periods of 215.2, 271.0 and 304.4 s are due to dipolar modes with $k=2, 3, 4$ respectively.  Our calculations  indicate that  the best fit  to the  observed period pattern of G117B-15A corresponds to  a DA white dwarf structure with a stellar  mass   of  0.525  \\msun,   with  a  hydrogen   mass  fraction \\lmh$\\gtrsim$-3.83  at an effective  temperature \\teff$\\approx$11800K. The  value  of the  stellar  mass  is  consistent with  that  obtained spectroscopically by Koester \\& Allard. ",
        "introduction": "\\label{sec:intro}  Asteroseismology  has  become a  powerful  method  to disentangle  the internal structure  and evolution  of stars by  means of the  study of their oscillatory pattern.  This  technique, very sophisticated in the case  of the Sun,  has also  undergone a  strong development  in other stars,  in  particular  in   variable  white  dwarf  (WD)  stars  (see e.g. Brown \\& Gilliland 1994; Gautschy \\& Saio 1995; 1996).  Pulsating WDs show multiperiodic luminosity variations in three ranges of  effective  temperatures  (\\teff)  corresponding to  the  currently called DOV,  DBV and DAV (see,  e.g., the review by  Winget 1988).  Of interest in this work are the DAVs (hydrogen - dominated atmospheres), or ZZ  Ceti, that  pulsate in the  instability strip  corresponding to 11000 K $\\leq$  \\teff $\\leq$ 13000 K.  The  periodicities in the light curves of pulsating WDs are  naturally explained in terms of nonradial g-modes  of  low  harmonic  degree  ($\\ell \\leq  2$),  driven  by  the ``$\\kappa$ mechanism'' working in  a partial ionization region near to the stellar  surface (Dolez  \\& Vauclair 1981,  Winget et  al.  1982). The periods ($P$) are found within a  range of $100 s \\leq P \\leq 1200 s$ and photometric amplitudes reach up to 0.30 magnitudes.  Most  asteroseismological studies  performed on  WDs to  date  rely on stellar models  constructed under some simplifying  hypotheses. One of the most relevant  is related to the profile  of the internal chemical composition in the interface  zones.  In such regions, the equilibrium diffusion in the trace  element approximation has been widely employed to  infer  the profile  of  the  chemical  distribution (see  Tassoul, Fontaine \\&  Winget 1990).  The  main motivation for  considering this approximation  is to  avoid  the solution  of  time dependent  element diffusion as  the WD evolves. In  the frame of  such an approximation, the profile  of the  interface region is  very simple:  its functional form is a power law. The  transition zone is separated into two parts: an  upper one  in  which one  element  is dominant  and  the other  is considered as  a trace  and a lower  region in  which the role  of the respective  elements  is  reversed.    Because  these  two  power  law solutions are matched for  fulfilling the condition of conservation of mass of each element, a discontinuity in the derivative occurs just at the matching point. The exponent of the power law solution is directly related to the  state of ionization of the  stellar plasma.  Thus, the structure of the  whole interface zone, in the  frame of this standard treatment,  can  be  modified  only  if  the  plasma  suffers  from  a modification  in  the state  of  ionization  as  a result  of  stellar evolution.  Calculations  of asteroseismology of  WDs in the  frame of such a  standard treatment for  the chemical interfaces are  those of, e.g., Bradley  (1996, 1998, 2001), Bradley \\&  Winget (1994), Brassard et al.  (1991,  1992ab), Fontaine et al. (1992),  Montgomery \\& Winget (1999),  Metcalfe,  Nather  \\&  Winget  (2000),  Metcalfe,  Winget  \\& Charbonneau (2001), Montgomery, Metcalfe \\& Winget (2001).  From  an  evolutionary  point  of  view, the  shape  of  the  chemical interfaces  may  not  be  critical\\footnote{An obvious  and  important exception is when the tail of chemical profiles are subject to nuclear burning.}, but in pulsational studies  they provide a non - negligible contribution   to   the   shape    of   the   Ledoux   term   of   the Brunt-V\\\"ais\\\"al\\\"a frequency (Brassard et  al.  1991).  Thus, we can, in  principle,  expect differences  between  studies with  equilibrium diffusion and those in which  other more physically sound treatment is performed. This  is a  very important point  to be made  in connection with the aim of the present work.  Since the  pioneering work of Iben  \\& MacDonald (1985),  we know that element diffusion modifies  the chemical abundance distribution within a WD star even during evolutionary stages corresponding to the ZZ Ceti domain (see Iben  \\& MacDonald 1985, particularly their  Fig.  4). Few calculations exist in the literature  in which the evolution of WDs is addressed  in  a  self  consistent  way with  time  dependent  element diffusion. Amongst them, we mention the study of MacDonald, Hernanz \\& Jos\\'e  (1998)  aimed  at   studying  the  carbon  pollution  in  cool WDs. Also, Dehner  \\& Kawaler (1995) used non  - equilibrium diffusion together  with  evolutionary calculations  to  study  WDs with  helium envelopes.   Finally,  Althaus,  Serenelli  \\& Benvenuto  (2001)  have recently shown that diffusion  induces the occurrence of thermonuclear flashes in  helium - core WDs,  causing the evolution of  such kind of WDs to  occur on timescales significatively shorter  than predicted by models  without diffusion.   This has  been particularly  important in solving the discrepancy  (Van Kerkwijk et al.  2000)  about the age of binary systems containing a millisecond pulsar and a helium WD.  As far  as we  are aware, the  only work  aimed at exploring  the role played by diffusion in the period and its rate of change for $g$-modes in  pulsating  DA  WDs is  that  of  C\\'orsico  et al.   (2001b).   In particular, the authors found that the differences in the shape of the chemical profiles at the interface zones induce appreciable changes in the periods, as  compared to the case of  equilibrium diffusion in the trace element  approximation.  Also,  there are noticeable  changes in the period derivative,  which are due in part to  the evolution of the chemical  profile during  the cooling  of the  WD across  the  ZZ Ceti instability strip.  Since  sometime ago  there have  been several  works available  in the literature  in which  the observed  period structure  of  a particular object is fitted to theoretical  predictions. Such works show that, in principle, information about the stellar mass and the stratified outer layer  structure can  be inferred  (see,  e.g., Winget  et al.   1991; Fontaine et al. 1992; Pfeiffer et al.  1996; Bradley \\& Kleinman 1997; Bradley 1998,  2001; Bradley  \\& Winget 1994).   As these  studies are based on  a simplified treatment of  diffusion, we believe  that it is worth revisiting this problem on  the basis of more detailed models of WD structure. In particular, in view of the results found in C\\'orsico et al. (2001b), we expect to find differences in the period fitting to a  particular object,  as compared  with  the situation  in which  the standard treatment is used. The only  way to find how important can be such differences is by performing a detailed asteroseismological study choosing a well studied object.  We consider G117-B15A  as an optimal target for  our study.  G117-B15A is an otherwise typical DA WD, the variability of which was discovered by Mc Graw  \\& Robinson (1976) and, since then,  it has been monitored continuously.  The mass and in particular the effective temperature of this   star  have   been   the  subject   of  numerous   spectroscopic redeterminations (see Gautschy, Ludwig \\& Freytag 1996 for a summary), in particular  values of  0.59 \\msun and  11620 K,  respectively, have been derived  by Bergeron et  al.  (1995).  More recently,  Koester \\& Allard (2000) (hereafter KA) have suggested a somewhat lower value for the  mass  of  0.53  \\msun  and  a  higher  effective  temperature  of \\teff=$11900 \\pm 140$ K.  G117-B15A has periods of oscillation of 215.2, 271 and 304.4 s (Kepler et al.  1982).  Notably, for the  215.2 s mode it has been possible to find a  value of its temporal  derivative (Kepler et al.   2000) to be $\\dot{P}= (2.3  \\pm 1.4) \\times 10^{-15}$  s s$^{-1}$.  Interestingly, the 215.2  s mode  present in G117-B15A  is the most  stable optically registered oscillation with a stability comparable to that of the most stable millisecond  pulsars.  As G117-B15A is a  well know oscillator, it has motivated  the interest of several researchers  in its internal structure.  For example, Bradley (1998)  has found the best fitting to the period pattern  with a model of $\\approx  0.6$ \\msun and depending on the  identification of the modes  (see below for  an explanation of its  meaning),  the favoured  value  for  the  hydrogen envelope  mass fraction \\mh is $-7  \\lesssim \\log{{\\rm (M_H/M_*)}}\\lesssim -6$ or $-5 \\lesssim  \\log{{\\rm  (M_H/M_*)}}\\lesssim   -4$,  and  \\lmhe=-2.   More recently, C\\'orsico  et al.   (2001a) have found,  on the basis  of WD models in which diffusion is neglected, the best fitting to the period pattern with  a 0.55 \\msun WD  model with a carbon  - oxygen interior, \\lmhe=-2 and \\lmh=-4 as  predicted by stellar evolution. Notably, such a fitting is  in nice agreement with one of  those proposed by Bradley (1998).  It  is the  aim  of  this work  to  perform a  fitting  to the  period structure present in G117-B15A  by computing the non-radial eigenmodes in the  frame of linear and adiabatic  approximation, and evolutionary WD  models  in which  time  dependent  element  diffusion is  properly accounted for.   About the way we  shall perform such  a fitting, some words are in order.  In handling  models like those we shall employ it is quite  obvious that we are  not in a position  to employ techniques like  those  used  in the  so  -  called  genetic algorithm  (see  its application for the  case of the DBV GD358 in  Metcalfe et al.  2000). In the  context of ZZ Ceti stars  a similar approach though  to a less extent has  been repeatedly  employed by numerous  investigations (e.g the complete and detailed study  performed by Bradley 1998, 2001).  Up to  now,  the  shape  and  thickness  of  the  relevant  chemical interfaces are usually  treated as a free parameter.   However, a more physically sound  treatment can be  performed when account is  made of time dependent  element diffusion in evolutionary models  of WDs. This aspect of WD evolution is one of the most important when an attempt is made  to  compare  observations  with  theoretical  expectations  from pulsating WDs.  Unfortunately,  there are currently some uncertainties (e.g., the treatment  of convection, the rate of  the critical nuclear reaction ${^{12}}{\\rm C}(\\alpha,\\gamma){^{16}}{\\rm O}$, wind mass loss episodes,  etc.) that  prevent us,  provided a  mass value  and pre-WD composition, from predicting a whole definite internal structure. This is especially  true for \\mh.  Thus,  in view of these  facts, we shall consider the  value of  \\mh as a  free parameter. We  assumed \\lmhe=-2 throughout this paper.  The  remainder  of the  paper  is  organized  as follows:  In  Section \\ref{sec:comput} we discuss the  physical ingredients we have employed together  with the  computational strategy  we have  employed. Section \\ref{sec:evolution}  is  devoted  to  presenting the  details  of  the evolutionary  models  we  have  constructed  down  to  the  conditions relevant   for   the   WD    we   are   investigating.    In   section \\ref{sec:asteroseis}  we present  the asteroseismological  analysis we performed   for   the   particular   case   of   G117B-15A.    Section \\ref{sec:dicuss} is devoted  to the discussion of the  results we have found and finally in Section \\ref{sec:conclus} we give some concluding remarks and insights for future investigations. ",
        "conclusions": "\\label{sec:conclus}  In this  paper we  have studied the  structural characteristic  of the variable  DA white  dwarf (WD)  G117B-15A by  applying the  methods of asteroseismology.  In  doing  so  we  have constructed  models  of  WD evolution considering  updated and detailed  physical ingredients.  It should be remarked than we have included several processes responsible for the diffusion  of elements in the WD  interior.  In particular, we considered  gravitational settling,  chemical  and thermal  diffusion. Starting  from  an  artificial   model,  we  have  considered  several thickness  for the  outermost hydrogen  layer, whereas  for  the inner helium-,  carbon- and  oxygen-rich  layer we  considered the  profiles predicted by  Salaris et al.  (1997).  The range of  stellar masses we have considered is $0.50 \\leq M_{*}/M_{\\odot} \\leq 0.60$. As far as we are aware, this is the first  study in which evolutionary models of DA WDs considering element  diffusion have been constructed for  a set of values of hydrogen mass fraction.  The evolution of each model  sequences were followed down to effective temperature  of 12500K  from  where on,  we  considered the  evolution coupled to the oscillations.  We considered dipolar, adiabatic g-modes with radial order  $k=1, \\cdots, 4$.  After constructing  the full set of  models, we considered  the location  of the  minima of  a function conveniently  defined for the  purpose of  fitting. We  considered two possible  mode  identifications  according  to  Bradley  (1998):  that observed  modes 215.2,  271 and  304.4 s  correspond to  $k=1,2,3$ and $k=2,3,4$.  We  find that the employment of  asteroseismology does not provide an  univocal answer about the correct  identification of modes present  in G117B-15A.  However, if  we  use them  together with  data deduced from spectroscopy  it is found that the  identification of the observed periods  of 215.2,  271 and 304.4  s with dipolar  modes with $k=2,3,4$ is clearly the best.  The results  presented in this  work (summarised in table  2) strongly suggest that G117B-15A is a DA  WD with a stellar mass of 0.525 \\msun, a  hydrogen   mass  fraction  \\lmh=-3.83   and  effective  temperature \\teff$\\approx$11800K. Notably, the values of the effective temperature and stellar  mass are in very  nice agreement with  those predicted by spectroscopic analysis  by Koester  \\& Allard (2000).  This represents the main result of the present paper.  While  the favoured value  for the  mass fraction  of hydrogen  is the maximum we  have considered, we  have not tried  to study the  case of slightly  (25\\% say)  thicker hydrogen  layers because  our artificial starting evolutionary technique prevents  us from being confident with the  possibility that  such structures  can be  actually be  formed in Nature.   The only  way to  answer this  question is  to  perform full evolutionary  calculations  from   the  initial  stages  of  evolution (hydrogen main sequence) to the WD stage.  Thus, performing full evolutionary studies coupled to oscillations is, in  principle,  interesting in  two  senses.   First,  it can  provide significative improvements to the  fitting of the oscillatory modes we have found for G117B-15A in this paper.  Second, and more importantly, it can  provide us with  a stringent test  of the quality  of detailed theoretical  evolutionary   models  of  WD  stars   by  comparing  the theoretical  pattern  of   non-radial  oscillations  against  accurate observational determinations  of the oscillatory pattern  of WD stars. Work in this sense is in progress and will be published elsewhere."
    },
    "0201/astro-ph0201010_arXiv.txt": {
        "abstract": "We present the results from our deep optical imaging survey \\mbox{$\\mu^{\\rm V}_{\\rm lim}\\!\\approx\\!26-27$\\,\\magsqarcsec} of a morphologically selected sample of 72 edge-on disk galaxies. The question of the global structure of galactic stellar disks, especially the radial surface brightness profile at large galactocentric distances, is addressed. We find that typical radial profiles are better described by a two-slope exponential profile ---characterised by an inner and outer scalelength separated at a break radius--- rather than a sharply-truncated exponential model. Results are given for three face-on equivalents, serving as the crucial test to assure the findings for the edge-on sample without possible geometrical line-of-sight effects. ",
        "introduction": "Van der Kruit (1979) initially found that the outer parts of disks of spiral galaxies do not retain their exponential light distribution to the observed faint levels, but rather show sharp edges. For three nearby edge-on galaxies (NGC\\,4244,\\index{object, NGC 4244} NGC\\,4565,\\index{object, NGC 4565} NGC\\,5907\\index{object, NGC 5907}) he derived that the typical radial scalelength $h$ steepens from 5\\,kpc to about 1.6\\,kpc at the edge of the disk. The existence of these truncations, which are already visible in contour maps of edge-on and even of some face-on galaxies, is now well accepted (Pohlen 2001), but no unique physical interpretation is given to describe this observational phenomenon. The proposed explanations span a rather wide range of possibilities. Van der Kruit (1987) deduced a connection to the galaxy formation process describing the truncations as remnants from the early collapse. Ferguson \\& Clark (2001), for example, proposed an evolutionary scenario represented by the viscous disk evolution models. And Kennicutt (1989) suggested a ---probably less striking--- star-formation threshold. Up to now the applied characteristic parameter for comparing the observational results of different studies is the distance independent ratio of a truncation radius $R_{\\rm t}$ to a measured radial scalelength $h$. Van der Kruit \\& Searle (1982) found for their sample of seven galaxies a value of $R_{\\rm t}/h\\eq4.2 \\pm 0.6$, whereas Pohlen, Dettmar, \\& L\\\"utticke (2000a) derived a significantly smaller one of $R_{\\rm t}/h\\eq2.9 \\pm 0.7$ for their CCD survey of 30 galaxies. ",
        "conclusions": ""
    },
    "0201/astro-ph0201360_arXiv.txt": {
        "abstract": "We present initial results of the first self-consistent numerical model of the outer magnetosphere of a pulsar. By using the relativistic ``particle-in-cell'' method with special boundary conditions to represent plasma dynamics in 3D, we are able to follow magnetospheric plasma through the light cylinder into the wind zone for arbitrary magnetic inclination angles. For aligned rotators we confirm the ``disk-dome'' charge-separated structure of the magnetosphere and find that this configuration is unstable to a 3D nonaxisymmetric diocotron instability. This instability allows plasma to move across the field lines and approach the corotating Goldreich-Julian solution within several rotation periods. For oblique rotators formation of the spiral ``striped wind'' in the equatorial direction is demonstrated and the acceleration of the wind and its magnetization is discussed. We find that the wind properties vary with stellar latitude; however, whether injection conditions at the pulsar are responsible for the observed jet-equator geometry of Crab and Vela is currently under investigation. We also comment on the electrodynamics of the simulated magnetospheres, their current closure, and future simulations. ",
        "introduction": "Most of the pulsar wind models assume a relativistic MHD-type flow with various degrees of symmetry (e.g., Kennel and Coroniti 1984, Begelman and Li 1992), and are reasonably successful in explaining the overall morphology and energetics of the surrounding plerion.  However, they run into various difficulties under attempts to extrapolate the solution back towards the pulsar. The most famous of these is the problem of the magnetization of the wind ($\\sigma$-paradox), which is the contradiction between the low ratio of magnetic to kinetic energy ($\\sigma \\sim 10^{-3}$) at the wind-nebula interface inferred from observations, {\\it expected} strong magnetization near the pulsar ($\\sigma \\gg 1$), and the conservation of $\\sigma$ in an ideal MHD flow. Recent observations of the Crab and Vela with Chandra and HST have underscored the importance of understanding the origins of the pulsar wind. Observations of pulsar jets and orthogonal equatorial outflows suggest that rotation of the neutron star is imprinted in the wind structure, possibly at a very early stage. How and where acceleration and collimation of the wind occurs can be reliably studied only by constructing a wind model that begins from the pulsar itself.  The main difficulty in the way of theoretical progress on this front is the lack of intuition about the phenomena occurring near a pulsar. For a general case of an oblique rotator, and, as we show in this paper, even for an aligned rotator, the problem is intrinsically three dimensional, nonaxisymmetric, and time-dependent, or in other words ``too complicated''. With observations unlikely to resolve light cylinder scales ($\\sim 10^8$ cm) any time soon, our only resort is numerical simulations. In this paper we describe our effort to simulate the formation of pulsar winds from the first principles. In particular, we study the behavior of plasma that is either emitted from or injected in the vicinity of a strongly magnetized rotating conducting sphere with arbitrary inclination angles between the magnetic and rotational axes. Despite its implied simplicity, this problem has very interesting solutions that may at first appear strange and unusual. Hence, in describing these preliminary results our emphasis is more on developing intuition, rather than on direct relevance to pulsars and observational implications. Such issues as well as more realistic models are currently being investigated. The paper is organized as follows: in \\S~2 we describe the particle-in-cell method used in our simulations and the setup of the problem, in \\S~{3} and \\S~{4} we describe numerical experiments with aligned and oblique rotators, and in \\S~{5} we conclude and discuss further work. ",
        "conclusions": "\\label{secconcl} Our findings can be summarized as follows: pulsar magnetospheres {\\it can} be filled by surface emission with the help of a 3D diocotron instability; simple models of oblique rotators {\\it can} form relativistic winds; such winds have variation in properties with latitude, with equatorial flow being formed at the edge of the closed magnetosphere, and polar cap flow probably contributing in the direction of rotational axis. Much work remains to be done to quantify these qualitative results and to understand the dependence of the wind geometry, acceleration and collimation on the parameter regime. While tempting, it would be too early to suggest that the equatorial and polar outflows observed in these simulations are directly responsible for features in the Crab and Vela. The models described in this paper are self-consistent in the sense that the field and particle equations are solved simultaneously and correctly. However, the models are not currently self-consistent with respect to the boundary conditions: the induced quadrupole electric fields are those from a rotating magnetized {\\it conductor}, but as far as the plasma is concerned the central body is a {\\it dielectric} with $\\epsilon=1$. Both KPM1, KPM2 and SMT simulations suffer from the same problem as well. This restriction does not allow currents to close inside the star, and instead leads to charging of the central body at the locations where the current is extracted and returned. Consequently, any current flow in the system is temporary. We are currently working on overcoming this deficiency. The more realistic models with true central conductors will be described in detail in an upcoming publication."
    },
    "0201/astro-ph0201156_arXiv.txt": {
        "abstract": "There is now strong evidence that many LLAGNs contain accreting massive black holes and that the nuclear radio emission is dominated by parsec-scale jets launched by these black holes. Here, we present preliminary results on the 1.4~GHz to 667~GHz spectral shape of a well-defined sample of 16 LLAGNs. The LLAGNs have a falling spectrum at high GHz frequencies. Several also show a low-frequency turnover with a peak in the 1-20~GHz range. The results provide further support for jet dominance of the core radio emission. The LLAGNs show intriguing similarities with gigahertz-peaked spectrum (GPS) sources. ",
        "introduction": "Low luminosity active galactic nuclei (LLAGNs), operationally defined as AGNs with nuclear H$\\alpha$ luminosity $<$~10$^{40}$ erg~s$^{-1}$, make up almost 50\\% of all nearby bright galaxies (e.g., Ho, Filippenko, \\& Sargent 1997a). They are spectroscopically sub-classified  into Low Ionization Nuclear Emission Region nuclei (LINERs), low luminosity Seyferts, and ``transition nuclei'', whose spectra are intermediate between Seyfert/LINER and HII region spectra.  Evidence has been accumulating that some fraction of LLAGNs share characteristics in common with more powerful AGNs. These similarities include the presence of nuclear compact radio cores (Heckman 1980), water vapor megamasers (Braatz et al. 1997), nuclear point-like UV sources (Maoz et al. 1995; Barth et al. 1998), broad H$\\alpha$ lines (Ho et al. 1997b), and broader H$\\alpha$ lines in polarized emission than in total emission (Barth et al. 1999). If LLAGNs are truly scaled down AGNs then the challenge is to explain their much lower accretion luminosities. This requires either very low accretion rates ($\\sim$10$^{-8}\\,$L$_{\\rm Edd}$; e.g. Falcke \\& Biermann 1999) or radiative efficiencies (the ratio of radiated energy to accreted mass) much lower than the typical value of $\\sim$10\\% (e.g.~Chapter~7.8 of Frank, King, \\& Raine 1995) assumed for powerful AGNs.  Our radio surveys of the 96 nearest LLAGNs from the Palomar spectroscopic survey  (Ho et al. 1997a) have found compact (150~mas) flat-spectrum radio cores in almost half of all LINERs and low-luminosity Seyferts (Nagar et al. 2000; 2002). Follow-up observations detected parsec-scale radio cores in all (16) LLAGNs with S$^{\\rm VLA}_{\\rm 2cm}$~$\\geq$ 2.7~mJy, implying brightness temperatures $\\ga$ 10$^8$~K (Falcke et al. 2000; Nagar et al. 2002). The five nuclei with the highest core fluxes - NGC~3031 (Bietenholz, Bartel, \\& Rupen 2000), NGC~4278 (Jones, Wrobel, \\& Shaffer 1984; Falcke et al. 2000), NGC~4486 (M~87; Junor \\& Biretta 1995), NGC~4374 (M~84; Wrobel, Walker, \\& Bridle 1996; Nagar et al. 2002), and NGC~4552 (M~89; Nagar et al. 2002) - all have pc-scale radio extensions, morphologically reminiscent of jets (e.g. Fig.~1). The high brightness-temperatures rule out an origin for the radio emission in star-formation processes. Thermal emission can also be ruled out as it would imply two to four orders of magnitude higher soft X-ray fluxes than seen in LLAGNs (see Falcke et al. 2000).  We are currently working on a high-resolution multifrequency radio study of all 16 LLAGNs at D $<$ 19~Mpc with confirmed pc-scale radio cores. Results on the 5~GHz to 15~GHz radio spectra are published in Nagar, Wilson, \\& Falcke (2001). Here, for the first time, we present preliminary results on the 1.4~GHz to 667~GHz spectral shape.  \\begin{figure}[t] \\plottwo{N4374.WOLABEL.PS}{N4552.WOLABEL.PS} \\caption{5~GHz VLBA maps of NGC~4374 (left) and NGC~4552 (right). The contours are integer powers of $\\sqrt{2}$, multiplied by the $\\sim3\\,\\sigma$ noise level of 1.2~mJy for NGC~4374, and by the $\\sim2\\,\\sigma$ noise level of 0.8~mJy for NGC~4552. The peak flux-densities are 152.6~mJy/beam and 93.8~mJy/beam, respectively. From Nagar et al. (2002).} \\end{figure} ",
        "conclusions": "First, we list the main caveats: (a)~as explained above the 5--15~GHz and the 22--43~GHz data have been normalized by the repeat 8.4~GHz observations. This assumes that the 5--43~GHz spectral shape did not change between the two epochs; (b)~the 1.4~GHz, 1.6~GHz, and the 353-667~GHz data are non-simultaneous with the other data; (c)~the resolution varies between 0{\\farcs}05 for the 43~GHz data and 13{\\arcsec} for the JCMT/SCUBA data. However, as discussed above, the radio cores are highly compact until mas-scales, so resolution effects should not cause significant bias at $\\geq$50~mas resolutions.  Despite these caveats, several deductions can be made from the results. First, let us restrict ourselves to scenarios in which the radio core is highly compact, e.g. advection dominated (ADAF) or convection (CDAF) dominated accretion flows (Narayan, Igumenshchev, \\& Abramowicz 2000). In these models, the radio emission is posited to come from the inner 10$^2$-10$^3$ Schwarzschild radii, i.e. highly sub-parsec for the black hole masses considered here. Thus, the radio fluxes we measure at all of our frequencies represent upper limits to the radio emission from such models. In fact in cases like this, it is not relevant to determine spectral indices from matched resolution radio maps; all that accomplishes is a deterioration of the few high resolution datapoints. We see no evidence of a moderately to highly inverted radio spectrum, as predicted by most ADAF and CDAF models (e.g. Quataert \\& Narayan 1999; Di~Matteo et al. 2001; summarized in Nagar et al. 2001). This tells us that even if such accretion flows exist, the sub-arcsecond radio emission upto 43~GHz (possibly all the way to 667~GHz) is dominated by other components.  On the other hand there is significant evidence for jets dominating the sub-arcsec radio emission in LLAGNs. The 5 nuclei with the highest core mas-scale fluxes all show extended emission, reminiscent of parsec-scale jets (see Section 1). In the three best studied nuclei - M~87, NGC~3031, and NGC~4258 - the `jet' dominates the `core' radio emission (see Nagar et al. 2001 for  a discussion of this). The radio spectral shape is also consistent with jet models: the high frequency spectral shape is consistent with optically-thin synchrotron emission, and the low frequency turnover could be caused by synchrotron self-absorption or by free-free absorption.  It is suggestive that the LLAGNs in Fig.~2 show several properties in common with gigahertz-peaked spectrum sources (GPSs; see O'Dea 1998 for a review): (a)~there is some evidence for a peak in the 1--20~GHz range; (b)~milli-arcsec imaging often shows a 'core-jet' or symmetric parsec-scale jets; (c)~LLAGNs in Fig.~2 have sub-pc scale radio jet extents and spectral peaks in the 5-10~GHz range. They thus extend to higher frequencies the relationship between linear size and turnover frequency in GPSs and Compact Steep Spectrum Sources (CSSs; see O'Dea et al. 1998). (d)~LLAGNs have a low X-ray luminosity (e.g. Ho et al. 2001), similar to the case of GPS galaxies.        \\newline In GPS sources, the radio emission is posited to come from the expanding lobes of the pc-scale radio jets. The low frequency turnover is deduced to be the result of synchrotron self-absorption or free-free absorption."
    },
    "0201/astro-ph0201406_arXiv.txt": {
        "abstract": "We have measured the distribution of intrinsic ellipticities for a sample of 28 relatively face--on spiral disks. We combine H$\\alpha$ velocity fields and $R$ and $I$-band images to determine differences between kinematic and photometric inclination and position angles, from which we estimate intrinsic ellipticities of galaxy disks. Our findings suggest disks have a log-normal distribution of ellipticities ($\\overline{\\epsilon} =0.06$) and span a range from $\\epsilon= 0$ (circular) to $\\epsilon=0.2$. We are also able to construct a tight Tully-Fisher relation for our face-on sample. We use this to assess the contribution of disk ellipticity on the observed Tully-Fisher scatter. ",
        "introduction": "Binney (1978) showed triaxial halos could affect disks by inducing warping and twists.  In particular, the axis ratio of halos lead to disks which are intrinsically elliptical (Franx \\& de Zeeuw 1992; Jog 2000). Hence the ellipticity of disks may plausibly be used to estimate the axis ratios of dark matter halos.  However, the inability to disentangle the ellipticity from the phase angle of this distortion makes such measurements difficult (e.g.  Zaritsky \\& Rix 1995; Schoenmakers 1999).  Andersen et al. (2001) presented a method which removed this degeneracy and yielded unique solutions for the disk ellipticity of nearly face-on galaxies by comparing kinematic and photometric inclination and position angles. This method assumes differences between these angles are solely the effect of ellipticity and not some other distortion.  Here we present results of a larger study to define the distribution of disk ellipticities and to establish if ellipticity is related to other physical quantities, e.g., Tully-Fisher (TF) scatter.  A sample of 39 galaxies were selected from the Principal Galaxy Catalog (Paturel et al. 1997) which have (1) {\\sl t}-types between 1.5--8.5, (2) axis ratios close to unity, (3) apparent disk sizes commensurate with the field of view of the integral field unit, DensePak, on the WIYN 3.5m telescope (Barden, Sawyer \\& Honeycutt 1998), and (4) low galactic absorption. We also required galaxies in the sample to be unbarred, isolated and have constant photometric position angles at three scale lengths.  \\begin{figure}[t] \\vbox to 2.15in{\\rule{0pt}{12in}} \\special{psfile=figure1.eps voffset=-150 hoffset=-27 vscale=65 hscale=65} \\caption[]{{\\bf Left Panel:} Differences in projected velocity as a function of $\\phi$ (the angle from the major axis in the {\\it observer's} frame) between rotating disks with inclination differences of 5$^\\circ$. These differences assume (1) measurements are made on the flat part of the rotation curve, (2) $V_{\\rm rot}=160 \\sin i$ km/s in the mean, and (3) orbits are circular. The solid curves represent mean inclinations of $15^\\circ, 25^\\circ, 35^\\circ, 45^\\circ, 55^\\circ, 65^\\circ$ and $75^\\circ$. The dashed line represents $\\theta=45^\\circ$ for each of these different inclinations, where $\\theta$ is the angle from the major axis in the {\\it galaxy} plane. Classical tilted-ring fits do not utilize data to right of dashed line, thereby missing over half the signal used to estimate inclination. {\\bf Right Panel:} The solid curve is our Monte Carlo prediction of inclination errors for velocity fields ``observed'' with two DensePak pointings and fit with a single inclined-disk velocity-field model, while points are errors measured using $\\chi^2$ intervals in fits to data. The dashed line represents $\\Delta i/i = 1$. Galaxies with $i>15^\\circ$ have inclination errors $\\Delta i< 5^\\circ$ which are sufficiently small to study the TF relation.} \\end{figure}  $R$ and $I$-band images were acquired at the WIYN 3.5m, KPNO 2.1m, McDonald Observatory 2.7m telescopes.  We used these images to measure axis ratios and position angles in a way designed to be unaffected by warps or spiral structure (see Andersen et al. 2001). H$\\alpha$ velocity fields were obtained using DensePak, feeding the WIYN Bench Spectrograph used with the 316 lines/mm echelle grating to cover 6500\\AA $< \\lambda\\lambda < 6900$\\AA, with an instrumental FWHM of 0.51 \\AA~(22.5 km/s).  Multiple DensePak pointings allowed us to map H$\\alpha$ velocity fields beyond 2.5 disk scale lengths. We modeled observed velocity fields to derive kinematic inclinations and position angles --- parameters critical to estimating disk ellipticity.  \\subsection{Velocity-Field Modeling}  Most galaxies in our sample do not show signs of rotation curve asymmetries, warps, solid body rotation, or spiral structure. Therefore, we adopted a single, inclined, differentially rotating, circular disk (``monolithic'') model to fit the DensePak H$\\alpha$ velocity fields instead of tilted ring models (e.g., Begeman 1989). There were two major advantages to our approach: (1) A monolithic velocity-field model uses all data to constrain the fit; and (2) a monolithic velocity-field model is better able to model low-inclination disks because tilted ring fits tend to diverge unless the fit is weighted by $|\\cos\\theta|$ ($\\theta$ is an angle measured from a galaxy's major axis in the galaxy plane) and data with $|\\theta|>\\theta_{max} = 45^\\circ$ is removed (Begeman 1989). However, the greatest differences between two velocity-field models with slightly different inclinations occur at $\\theta>45^\\circ$, precisely where tilted-ring fits often do not consider the data (left panel of Figure 1). Since differences between velocity fields of different inclination decrease with inclination, it is imperative to use data at all azimuthal angles to accurately fit velocity-field models at low inclinations, i.e. $i<30^\\circ$ (right panel of Figure 1).  A hyperbolic tangent rotation curve was sufficient to fit the shape of rotation curves in our sample with a minimum of free parameters. Our model had the following free variables: inclination, position angle, center, central velocity, observed rotation velocity, and hyperbolic tangent scale-length. The results of the model fits indicate our approximation that orbits are circular appears to be acceptable; for $\\epsilon_D<0.2$ the model inclination and position angles derived from circular versus elliptical orbits would be quite similar. We determined kinematic parameters for 36 of 39 galaxies using our fitting procedure.  Of the three galaxies for which we could not fit velocity-field models, two were at very low inclinations, while the third had insufficient data.  \\subsection{Results}  We derive ellipticities for the 28 of 39 galaxies for which accurate measures of the photometric and kinematic indices exist. We find a mean disk ellipticity of $\\overline{\\epsilon_D}=0.076$ If we assume the halo potential is non-rotating and has a constant elliptical distortion, we can estimate a halo ellipticity of $\\overline{\\epsilon_\\Phi}=0.054$ which is consistent with previous estimates of halo ellipticity (Rix \\& Zaritsky 1995; Schoenmakers 1999). Our unique solutions for disk ellipticity also allow us to determine the distribution of ellipticities for our sample, which we find is well-fit by a log-normal distribution with a mean and standard deviation on $\\ln(\\epsilon_D)$ equal to -2.82$\\pm0.73$ ($\\epsilon_D=0.060^{+0.064}_{-0.031}$, left panel of Figure 2).  \\begin{figure}[t] \\vbox to 1.8in{\\rule{0pt}{10in}} \\special{psfile=figure2.eps voffset=-150 hoffset=-20 vscale=55 hscale=55} \\caption[]{{\\bf Left Panel:} Distribution of disk ellipticities for our sample of 28 galaxies.  This distribution is well-fit by a log-normal distribution (dashed line) characterized by a mean ellipticity $\\overline{\\epsilon_D}=0.060$.  {\\bf Center Panel:} A Tully-Fisher relation for a sample of galaxies with a mean inclination of 26$^\\circ$. The dashed line represents the best fit TF relation to a subsample of galaxies in the quiet Hubble flow taken from Courteau (1997) Only 0.44 magnitudes of scatter was exhibited about this relation.  {\\bf Right Panel:} Component of TF scatter due to assuming circular orbits for an elliptical potential (Franx \\& de Zeeuw 1992; Table 1) versus TF scatter for our sample of nearly face-on galaxies. } \\end{figure} ",
        "conclusions": ""
    },
    "0201/astro-ph0201295_arXiv.txt": {
        "abstract": "The  Crab Nebula is likely to be expanding into freely expanding supernova ejecta, although the energy in the ejecta may be less than is typical for a Type II supernova. Pulsar nebulae much younger than the Crab have not been found and could have different properties. The search for such nebulae through ultraviolet/optical line emission in core collapse supernovae, or through their X-ray emission (which could show strong absorption) is warranted. Neutron stars have now been found in many young supernova remnants. There is not a clear link between neutron star and remnant type, although there is an indication that normal pulsars avoid the O-rich remnants. In the later phases of a supernova remnant, the pulsar wind nebula is crushed by the reverse shock front. Recent simulations show that this process is unstable, which can lead to mixing of the thermal and relativistic gases, and that the pulsar nebula is easily displaced from the pulsar, which can explain the position of the Vela pulsar relative to the Vela X radio nebula. ",
        "introduction": "Pulsars are thought to be born in the core collapse and explosion of massive stars, so that the initial evolution and observability of the PWNe (pulsar wind nebulae) depends on the properties of the supernova and the supernova remnant evolution. Several phases of evolution are expected (Pacini \\& Salvati 1973; Reynolds \\& Chevalier 1984; Chevalier 1998). Initially the radiation from the PWN is absorbed by the surrounding supernova gas; high energy photons may be reprocessed and appear at a lower energy. After $\\ga$10's of years, the supernova  becomes optically thin and the PWN radiation directly escapes. The Crab Nebula may be in this phase of evolution. After $\\sim 10^3$ years for a pulsar like that in the Crab, the spindown power from the pulsar decreases and the PWN fades. After $\\sim 10^4$ years, the reverse shock wave from the supernova remnant moves to the center of the remnant, crushing the PWN and leading to an increase in brightness of the nebula.  The discussion of PWNe has typically centered on objects like the Crab and its pulsar, which is estimated to have an initial spin period $P_0 \\approx 19$ ms and a magnetic field of $5\\times 10^{12}$ G. It is thought to belong to the main (not binary related) group of radio pulsars, which appear to have a fairly narrow distribution of initial magnetic field strengths $\\sim 3\\times 10^{12}$ G (Stollman 1987). The initial rotation rates of these pulsars remain uncertain; studies of the pulsar population are influenced by selection effects (e.g., Emmering \\& Chevalier 1989). In recent years there has been increasing evidence for a class of strongly magnetized neutron stars, or magnetars, with $B\\ga  10^{14}$ G (Thompson \\& Duncan 1995). Magnetars are thought to be associated with soft gamma-ray repeaters and possibly with AXPs (anomalous X-ray pulsars), both of which appear to contain slowly rotating neutron stars. Although the youth of AXPs is indicated by their association with supernova remnants, there is no evidence for PWNe. Finally there is a class of ``quiet'' neutron stars that have been found as weak X-ray sources in young supernova remnants, such as Cas A (Tananbaum 1999) and Puppis A (Petre, Becker,  \\& Winkler 1996). Although the relative formation rates of the normal radio pulsars, magnetars, and quiet neutron stars is not known, the likely youth of the objects in the latter two categories indicates that all three groups make a significant contribution to the total rate of massive star core collapses.  Here, I discuss four topics related to neutron stars and their supernova remnants. The first is the quest for young PWNe, considerably younger than the Crab Nebula. Next, I consider the status of the Crab Nebula and its surroundings. The third topic concerns the possible relation between a neutron star and its surrounding supernova, given the growing number of such associations. Finally, I review recent investigations of the effect of a reverse shock front on a PWN. ",
        "conclusions": "Of the various phases expected for pulsar nebula evolution, the earliest phase when nebular emission is absorbed by the surrounding supernova is least understood because of a lack of observations. The next phase when the PWN can be observed directly and is expanding into freely expanding supernova ejecta is presumably observed in many cases, although clear proof of this phase has been difficult to obtain even in the case of the Crab Nebula. Most of the remnants listed in Table 1, with ages $\\la 3000$ yr, are likely to be in this phase. Continued study of these sources, especially analysis of the composition from X-ray emission, is warranted in order to search for correlations between neutron stars and their progenitor stars. In the late, post-reverse shock phase, a pulsar nebula is easily distorted and displaced by the thermal gas. Careful studies of the boundaries of these nebulae should reveal the nature of the interaction."
    },
    "0201/astro-ph0201540_arXiv.txt": {
        "abstract": "{ We study luminosity and morphology segregation of cluster galaxies in an ensemble cluster built from 59 rich, nearby galaxy clusters observed in the ESO Nearby Cluster Survey (ENACS). The ensemble cluster contains 3056 member galaxies with positions, velocities and magnitudes; 96\\% of these also have galaxy types. From positions and velocities we identify galaxies within substructures, viz. as members of groups that are significantly colder than their parent cluster, or whose average velocity differs significantly from the mean. \\\\ We compare distributions of projected clustercentric distance $R$ and relative line-of-sight velocity $v$, of galaxy subsamples drawn from the ensemble cluster, to study various kinds of segregation, the significance of which is obtained from a 2-dimensional Kolmogorov-Smirnov test. We find that luminosity segregation is evident only for the ellipticals that are outside (i.e. not in) substructure and which are brighter than $M_R = -22.0\\pm 0.1$. This is mainly due to the brightest cluster members at rest at the centre of the cluster potential. \\\\ We confirm the well-known segregation of early- and late-type galaxies. For the galaxies with $M_R > -22.0$ of {\\em all} types (E, S0, S and emission-line galaxies, or ELG, for short), we find that those within substructure have $(R,v)$-distributions that differ from those of the galaxies that are not in substructure. The early and late spirals (Sa--Sb and Sbc--Ir respectively) that are not in substructure also appear to have different $(R,v)$-distributions. For these reasons we have studied the segregation properties of 10 galaxy subsamples: viz. E, S0, S$_e$, S$_l$ and ELG, both within substructure and outside substructure. \\\\ Among the 5 samples of galaxies that are {\\em not in substructure}, at least 3 ensembles can and must be distinguished; these are: [E+S0], S$_e$, and [S$_l$+ELG]. The [E+S0] ensemble is most centrally concentrated and has a fairly low velocity dispersion that hardly varies with radius. The [S$_l$+ELG] ensemble is least concentrated and has the highest velocity dispersion, which increases significantly towards the centre. The class of the S$_e$ galaxies is intermediate to the two ensembles. Its velocity dispersion is very similar to that of the [E+S0] galaxies in the outer regions but increases towards the centre. \\\\ The galaxies {\\em within substructure} do not all have identical $(R,v)$-distributions; we need to distinguish at least two ensembles, because the S0 and [S$_l$+ELG] galaxies have different distributions in $R$ as well as in $v$. The [S$_l$+ELG] galaxies are less centrally concentrated and, in the inner region, their velocity dispersion is higher than that of the S0 galaxies. Our data allow the other 3 galaxy classes to be combined with these two classes in 4 ways. \\\\ We discuss briefly how our data provide observational constraints for several processes inside clusters, like the destruction of substructure, the destruction of late spirals and the transformation of early spirals into S0's. ",
        "introduction": "\\label{s-intro}  It has been known for a long time that in clusters, galaxies of different classes have different projected distributions. Oemler (\\cite{oe74}), Melnick \\& Sargent (\\cite{ms77}) and Dressler (\\cite{dr80a}) were the first to quantify these differences. Dressler (\\cite{dr80a}) showed that the different distributions arise mainly from the so-called morphology-density relation (MDR): i.e., the relative fractions of ellipticals, S0's and spirals correlate very well with local surface density. Hence, the composition of the galaxy population changes with distance from the cluster centre.  Postman \\& Geller (\\cite{pg84}) derived an MDR over 6 decades of local space density from the CfA redshift survey, and in the Pisces-Perseus supercluster -- and in particular in its long filament -- Giovanelli et al. (\\cite{gi86}) found a clear MDR. In this supercluster, even early and late spirals have different distributions, and this was also found for spirals in groups of galaxies (Giuricin et al. \\cite{gi88}). The MDR was also studied in several individual nearby clusters (e.g.  Andreon \\cite{an94}, \\cite{an96}; Caon \\& Einasto \\cite{ce95}) and in general redshift surveys (e.g. Santiago \\& Strauss \\cite{ss92}).  In spite of the wealth of observational data, it is still not totally clear how the MDR arises. In clusters, galaxy encounters must play a r\\^ole, so that gas-rich disk galaxies cannot survive in the dense cores of clusters. Contrary to Dressler (\\cite{dr80a}), Whitmore \\& Gilmore (\\cite{wg91}) and Whitmore et al. (\\cite{wi93}) found that morphological fraction correlates as tightly with clustercentric distance as with projected density. An explanation for the MDR in a cold dark matter-dominated universe was given by Evrard et al. (\\cite{ev90}).  The study of the MDR was extended towards higher redshifts, e.g. by Dressler et al (\\cite{dr97}), Couch et al. (\\cite{co98}) and Fasano et al. (\\cite{fa00}), and was linked to the more general question of the evolution of galaxies in environments of different densities (e.g. by Menanteau et al. \\cite{me99}). Dressler et al. (\\cite{dr97}) found that the regular, centrally concentrated clusters at a redshift of about 0.5 show a strong MDR, as do the low-redshift clusters. However, the less concentrated and irregular clusters at $z \\approx$ 0.5 do not show a clear MDR, unlike their low-redshift counterparts.  Dressler et al. also noted that the fraction of S0's appears to decrease quite strongly with increasing redshift (by as much as a factor of 3 from $z = 0$ to $z \\approx 0.5$), and Fasano et al. (\\cite{fa00}) studied this effect in clusters at redshifts between 0.1 and 0.25.  The reality of this decrease was questioned by Andreon (\\cite{an98}) who argued that it is not trivial to establish a reliable elliptical/S0-ratio.  Also, the elliptical/S0-ratio may not be very meaningful (even if it can be established accurately) because the differences between ellipticals and S0's may not be major (e.g. J{\\o}rgensen \\& Franx \\cite{jf94}).  Morphological segregation in position is often accompanied by morphological segregation in velocity space, i.e.  galaxies of different types have different velocity dispersions, or velocity dispersion profiles (e.g. Tammann \\cite{ta72}; Moss \\& Dickens \\cite{md77}; Sodr\\'e et al. \\cite{so89}; Biviano et al. \\cite{bi92}). The effect is sometimes reported as a correlation between kinematics and colours (e.g. Colless \\& Dunn \\cite{cd96}; Carlberg et al. \\cite{ca97a}).  Luminosity segregation was detected by Rood \\& Turnrose (\\cite{rt68}), Capelato et al. (\\cite{ca80}), Yepes et al. (\\cite{ye91}) and Kashikawa et al. (\\cite{ks98}).  Luminosity segregation was detected both as a segregation in clustercentric distance, and as a kinematical segregation, viz. the most luminous galaxies have the smallest velocity dispersion (see e.g. Rood et al. \\cite{ro72}). Yet, kinematical segregation appears to occur mostly (Biviano et al. \\cite{bi92}), if not exclusively (Stein \\cite{st97}), for ellipticals, and much less -- if at all -- for the other galaxy types. Fusco-Femiano \\& Menci (\\cite{fm98}) explained the observed degrees of luminosity segregation by their merging models.  Adami et al. (\\cite{ad98a}) studied a sample of about 2000 galaxies in 40 nearby Abell clusters and confirmed that the overall velocity dispersion depends on galaxy type, and increases along the Hubble sequence. The velocity dispersion profiles for the various galaxy types indicate that the spirals may not yet be fully virialized, and may still be mostly on radial, infalling orbits. The spirals may thus have properties similar to the galaxies with emission lines (ELG), which were studied in A576 by Mohr et al. (\\cite{mo96}), and by Biviano et al. (\\cite{bi97}, hereafter Paper III) in the clusters observed in the ESO Nearby Abell Cluster Survey (ENACS). Biviano et al. concluded that the ELG probably have a significant velocity anisotropy. De Theije \\& Katgert (\\cite{tk99}, hereafter Paper VI) distinguished early- and late-type galaxies in the ENACS from their spectra, and concluded that the evidence for radial orbits was only significant for the ELG.  Recently, Thomas \\& Katgert (\\cite{tk02}, hereafter Paper VIII) derived morphologies for close to 2300 ENACS galaxies from CCD imaging. By adding morphologies from the literature, and spectral types from the ENACS spectra, this provides essentially complete type information for the galaxies in a sample of 59 ENACS clusters. This dataset allows a vastly improved analysis of the distribution and kinematics of the various classes of galaxies in clusters, which we present in this paper. In a subsequent paper (Katgert et al. \\cite{ka02}) we derive the mass profile in the ENACS clusters.  In Sect.~\\ref{s-data} we summarize the data that we used. In Sect.~\\ref{s-segr} we discuss the method by which we study the various types of segregation. In Sect~\\ref{s-subs} we discuss the effect of substructure and in Sects.~\\ref{s-lums} and ~\\ref{s-morphs} we discuss the evidence for luminosity and morphology segregation, as well as the minimum number of galaxy ensembles that must be distinguished. In Sect.~\\ref{s-nat_seg} we discuss the nature of the morphological segregations and in Sect.~\\ref{s-disc} we discuss the implications of our results for ideas about cluster galaxy evolution. In Sect.~\\ref{s-summ} we present a summary and the main conclusions. ",
        "conclusions": "\\label{s-summ}  We have studied evidence for luminosity and morphology segregation in an ensemble cluster of $\\sim 3000$ galaxies with positions, magnitudes, velocities, and galaxy type, in clusters observed in the ESO Nearby Abell Cluster Survey. From positions and velocities we identify galaxies in and outside substructure. The fraction of galaxies in substructure appears to decrease strongly towards the cluster centre. Luminosity segregation is evident only for the very bright ($M_R \\la -22$) ellipticals outside substructure, which mostly are brightest cluster members near the centres of their clusters.  For galaxies of all types, we find that those within substructure are segregated with respect to those outside substructure. This is mostly due to the fact that galaxies in and outside substructure have very different radial distributions. In addition, morphology segregation is found among galaxies both {\\em in} and {\\em outside} substructure.  The early- and late-type galaxies {\\em outside substructure} have different $(R,v)$-distributions, i.e. of projected position $R$ and relative velocity $v$. The early-type galaxies except the brightest ellipticals all have very similar $(R,v)$-distributions, i.e. the fainter ellipticals and S0's are not segregated. Similarly, the late spirals and the emission-line galaxies have indistinguishable $(R,v)$-distributions, but the $(R,v)$-distributions of the early spirals differs from that of the early-type galaxies and from that of the other late-type galaxies. Among galaxies {\\em in substructure}, the S0's are segregated from the late spirals and the emission-line galaxies, separately as well as together.  Luminosity segregation is most likely due to the dissipative processes in the innermost region of the clusters, which presumably produce the brightest ellipticals. The decrease of the fraction of galaxies in substructure towards the centre is probably due to tidal disruption. The cause for the accompanying decrease of the fraction of late-type galaxies in the subclumps is not evident, in particular because the latter is even stronger than the decrease towards the centre of the fraction of late-type galaxies outside substructure.  The large difference between the radial distributions of early and late spirals are attributed to systematic differences in their mass profiles. The late spirals presumably are fairly easily destroyed through impulsive encounters with other galaxies, and the early spirals much less, so that they can survive in the inner cluster regions. We briefly discuss the constraints that our data provide for the process by which early spirals transform into S0's."
    },
    "0201/astro-ph0201248_arXiv.txt": {
        "abstract": "We use the constraints arising from primordial nucleosynthesis to bound a putative electric charge density $|e|n_q$ of the universe. We find $|n_q/n_\\gamma| \\lesssim 10^{-43}$, four orders of magnitude more stringent than previous limits. We also work out the bounds on $n_q$ in models with a photon mass, that allows to have a charge density without large-scale electric fields. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0201/astro-ph0201138_arXiv.txt": {
        "abstract": "We report new radio continuum observations with an angular resolution of 2$''.1$ at 1.4 GHz (20 cm) and $0''.28$ at 8.4 GHz (3.6 cm), of the barred galaxy NGC 3367. In the map at 1.4 GHz, the central nuclear region connects to the SW lobe, with a projected structure of at a position angle of P.A.$\\sim230^{\\circ}$, forming a jet-like structure.  The map at 8.4 GHz shows a compact unresolved source (smaller than $65\\,$ pc in diameter) associated with emission from the nucleus and several compact sources located within a radius of about 300 pc, forming a circumnuclear structure. The compact core, jet, and lobes form a small, low-power counterpart to radio galaxies, with a flow axis that is out of the plane of the galaxy. The flow axis (P.A.$\\sim230^{\\circ}$) coincides with the P.A. of the major axis of the galaxy and is thus inclined to the rotation axis of the disk. In addition, the flow axis differs by about 20$^{\\circ}$ from the major axis of the stellar bar. Assuming that the stellar bar rotates counterclockwise (ie. assuming trailing spiral arms), this difference in angle is taken as an argument in favor of having the jet-like structure out of the plane of the disk and not associated with the stellar bar. ",
        "introduction": "Strong radio continuum emitters, such as radio galaxies or quasars, are identified with elliptical galaxies and recent mergers, and many of them have impressive radio jets with sizes larger than the host galaxy (see Wilson \\& Colbert 1995).  Maps of radio galaxies display fairly straight radio jets with distant lobes, showing the interaction of the jet with the ambient medium and indicating that the jet axis has been stable for thousands or even millions of years. Recent work indicates the presence of lobes at a distance of hundred kiloparsecs from a spiral galaxy in a cluster (\\cite{led98,led01}). In normal spiral galaxies, on the other hand, most of the radio continuum emission comes from the disk component (\\cite{hum81,con87,gar93,nik97}). Radio surveys of Seyfert galaxies indicate that the radio continuum emission arises mainly from three components: (1) subkiloparsec emission from the nuclear region; (2) extranuclear kpc scale emission; and (3) greater than kpc-scale emission associated with the disk (\\cite{wil83,ulv84,bau93,col96,ho01}). In barred spiral galaxies, the radio continuum emission is found from the following: (1) emission from the compact nucleus; (2) emission from circumnuclear region ($\\leq 1$kpc); (3) emission from dust lanes in leading side of the stellar bar; and (4) emission from spirals arms and disk (\\cite{hum87,con87,gar91a,gar91b,lin99,bec99}). Depending on the central activity, a spiral can be classified as a starburst or an active galaxy. Seyfert galaxies are stronger radio emitters at 1.4 GHz than normal and barred spirals (\\cite{con87}).  The central radio sources are sometimes associated with a pair of extended sources identified as lobes (\\cite{ulv84,ulv87,bau93,col96,ho01}).  These triple sources (compact nucleus plus lobes) are thought to be small scale, low power versions of the large scale jets and lobes seen in radio galaxies and quasars. The radio continuum emission from the central region of spirals is either linked to star formation, via HII regions and supernova remnants as in starbursts, or to an unresolved compact object, probably an AGN (Baum et al 1993). The study of NGC 3367 may be important in order to understand the lobe-central source relationship and help to clarify the starburst-AGN dichotomy.  The barred spiral, NGC 3367, is a mildly active galaxy, between a weak Liner and a HII nucleus (\\cite{ver97,ho97}) that displays such a triple source structure. At 15$''$ angular resolution, in addition to extended emission from different locations in the disk, it shows an unresolved radio source in the center plus two sources in opposite directions from it (\\cite{con90}). Maps at 4$''.5$ resolution indicated more clearly the presence of the two lobes, (\\cite{gar98}). The central source was not resolved at this 4$''.5$ angular resolution, and it was unclear if it is a point-like source or an extended circumnuclear structure, and whether or not it is really connected with the lobes.  In this paper we present new VLA\\footnote{The VLA is part of the National Radio Astronomy Observatory which is a facility of the National Science Foundation operated under agreement by Associated Universities Inc.} radio continuum observations of NGC~3367 at 1.4 GHz with a beam of $\\approx2''.1$, and at 8.4 GHz with a beam of $\\approx0''.28$. \\S 2 presents the observed properties of NGC 3367, \\S 3 presents the new radio continuum observations and results, and \\S 4 gives the discussion and conclusions.  \\section {NGC 3367}  NGC~3367 is an SBc(s) barred spiral galaxy with a stellar bar structure of diameter $\\approx 32''$ (6.7 kpc) oriented at a position angle (PA) $\\approx 70^{\\circ}$. The disk of NGC 3367 is inclined with respect to the plane of the sky at an angle between $6^{\\circ}$ (Grosb\\o l 1985) and $30^{\\circ}$ (\\cite{gar01}), and has an optically bright SW structure resembling a half-ring, or a large scale ``bow shock'', at about 10 kpc from the nucleus.  This structure is formed by a collection of H$\\alpha$ regions, that looks like a necklace and its origin has been ascribed to an off-center impact with an external intruder, most likely a small galaxy (\\cite{gar96a}).  Indeed, the general arrangement of the H$\\alpha$ knots is similar to the elongated rings found in numerical simulations of off-center galaxy collisions by Gerber \\& Lamb (1994).  The expanding density wave created by the collision can trigger the formation of the half-ring of HII regions and, given that the expected wave velocity in the disk is below 100 km/s, the collision probably ocurred a few times 10$^8$ yr ago.  In addition, aside from the radial gas inflows induced by the stellar bar, a galactic collision is also able to drive gas towards the galactic center inducing circumnuclear star formation as well as nuclear activity. NGC 3367 also shows H$\\alpha$ emission from the central region with an unresolved source, most likely, a combination of emission from a compact source and circumnuclear structure at a radius of less than 500 pc (\\cite{gar96a,gar96b}).  The disk of NGC 3367 has a normal content of atomic hydrogen, with M$_{HI}\\sim 7\\times10^9$ M$_{\\odot}$ (\\cite{huc85}), and it is considered an isolated field galaxy, behind the Leo group of galaxies at a distance of 43.6 Mpc, with its closest neighbor more than 900 kpc away to the NE (\\cite{tul88}).  Its optical spectrum shows moderately broad H$\\alpha$+[NII] lines with FWHM$\\sim650$ km s$^{-1}$, but H$\\beta$ is stronger than [OIII]$\\lambda 5007$\\AA~ and there is weak emission of He II $\\lambda 4686$\\AA, suggesting the existence of WR stars and weak Liner activity (\\cite{ver86,ver97,ho97}). In addition, its X~ray luminosity is stronger than that of any normal spiral galaxy, but weaker than Seyfert or radio galaxies (\\cite{gio90,sto91,fab92}).  Nonetheless, from the optical line ratios, NGC 3367 is also considered to have an HII nucleus (\\cite{ho97}).  The first radio continuum observation of the region around NGC 3367 was at 178 MHz, with several arcmin angular resolution (\\cite{gow67}). Although the radio emission was identified with NGC 3367 (\\cite{cas67}), the bulk of the emission most likely originated from a radio galaxy $\\approx3'$ north of NGC 3367, detected later with better angular resolution (\\cite{law83}). Similarly, observations from Arecibo at 430 MHz and 835 MHz with integrated fluxes of 583 mJy and 365 mJy respectively with $\\approx9'$ angular resolution most likely included the flux of the background radio galaxy (\\cite{isr83}). Green Bank single dish observations of NGC 3367 at 5 GHz with a resolution of $\\sim3'$ were carried out by Sramek (1975) and Bennett et al. (1986) with integrated fluxes of 35 mJy and 71 mJy respectively. Israel \\& van der Hulst (1983) reported an integrated flux of 18 mJy at 10.7 GHz with an angular resolution of $3'$ from OVRO 40m. Finally Dunne et al. (2000) reported a flux of 132 mJy at 350 GHz with $15''$ angular resolution.  The first aperture synthesis map of NGC 3367 at 1.49 GHz, with 15$''$ angular resolution, showed diffuse disk emission plus three peaks aligned in the NE -- SW direction (\\cite{con90}), and the disk emission was noticed to be edge brightened.  A more recent mapping at this same frequency, but now with 4$.''5$ angular resolution, shows more clearly the emission from the triple source: emission from the nuclear region in addition to emission from two extended lobes at a distance of $\\sim 6$ kpc from the center (\\cite{gar98}).  The polarization analysis indicated that only the SW lobe is polarized, suggesting that it is out of the disk of the galaxy and closer to the observer than the NE lobe (the emission of the NE lobe has been depolarized because this emission has passed through the plane of the galaxy)(\\cite{gar98}). These observations showed very clearly the presence of kiloparsec scale lobes from a barred spiral galaxy seen almost face on, in addition to the weaker emission from a large number of compact sources in the disk. ",
        "conclusions": "We have observed the radio continuum emission from the barred galaxy NGC~3367, with the VLA A array, at 1.4 GHz and 8.4 GHz with 2$''.1$ and $0''.28$ spatial resolutions, respectively.  The radio maps show emission from the central region, the lobes, and a weak jet. This morphology in NGC 3367 resembles the morphology of more powerful radio galaxies (\\cite{far74,sch01}). The power at 1.4 GHz is two to three orders of magnitude less than any radio galaxy and it is more like the power of Seyfert galaxies, see Table 1 (see Figure 3 of \\cite{ho01,ulv01}). Nonetheless this is one of the largest and best defined triple source ever detected in a galaxy considered to be a normal barred spiral. The triple source is due to emission from the central sources, emission from a jet-like structure connecting the central emission with the lobes and the extended lobes. Other large radio structures have been observed in the Seyfert galaxies , for example, Mrk 6 (lobe extent 14 kpc), Mrk 348 (lobe extent 5 kpc), NGC 3516 (lobe extent 8.5 kpc) (\\cite{bau93}), NGC 4235 (lobe extent 9 kpc) (\\cite{col96}) and in the disk galaxy O313-192 in the cluster A428 (lobe extent 100 kpc) (\\cite{led98,led01}). The radio continuum jets in other barred galaxies as NGC 1068 (\\cite{wil87}) and NGC 5728 (\\cite{sch88}) are much smaller in size (of the order of tens to hundred parsecs). In contrast with O313-192, which has an AGN, and the Markarian galaxies, which are Seyfert 1s and 2s, NGC 3367 is only a mildly active Liner and, based on optical line emission ratios and widths, is not even considered to be a Seyfert galaxy.  \\begin{figure}[t] \\psfig{figure=garcia3367.fig5.ps,width=8cm,angle=0} \\caption[figure=garcia3367.fig5.ps]{Radio continuum emission from the innermost central region at 1.4 GHz, in contours, and at 8.4 GHz, in greyscale. The contours are in units of 25$\\mu$Jy/beam and the levels are 1, 1.5, 2, 3, 6, 10, 30, 50, 100, 150, 250 and 390. The greyscale is from 30$\\mu$Jy/beam to 200$\\mu$Jy/beam.} \\label{fig5} \\label{hi} \\end{figure}   The P.A$._{jet}$ of the jet-like structure ($\\sim230^{\\circ}\\pm10^{\\circ}$; see Figure 3) is the same as the P.A$._{ma}$ of the major axis, as determined from H$\\alpha$ kinematics (\\cite{gar01}).  The jet would be considered a low power one (\\cite{mas96}). The direction of the plasma outflow is thus not aligned with the rotation axis of the disk because, if that were the case, the relative orientation of the lobes, in NGC 3367, would have to lie closer to the direction SE-NW, that is, the projection of the P.A. of the minor axis as expected in the very simple picture that jets emanating from an active nucleus would emerge at right angles to the disk of the host galaxies (\\cite{kin00}).  However the simple scenario is contradicted by the observations of Seyfert galaxies (\\cite{sch97,kin00,ulv01}). Our observations of NGC 3367 suggest that the outflow axis is inclined with respect to the rotation axis of the galaxy and also inclined to the line of sight. A comparison between the the P.A. of the extended radio structures from Seyferts with their host galaxies' major-axis P.A. indicate that the radio structures in type-2 Seyferts are oriented along any direction in the galaxy, and not necessarily along the minor axis (\\cite{sch97,kin00,ulv01}). The directions of the radio jets are consistent with being completely uncorrelated with the planes of the host galaxies (\\cite{pri99,nag99,kin00}). Our observations indicate that NGC 3367 (being a noninteracting late-type Liner/HII spiral having kpc lobes) presents P.A$_{radio}\\approx$P.A$_{maj}$.  \\begin{figure}[t] \\psfig{figure=garcia3367.fig6.ps,width=8cm,angle=0} \\caption[figure=garcia3367.fig6.ps]{H$\\alpha$ - 8.4 GHz emission from the innermost $4''$. The contours and greyscale are in arbitrary units relative to the maximum. The structure observed indeed suggests the existence of a circumnuclear structure within the innermost 300 pc ($1''=210$ pc). This image was obtained with several assumptions, among them are (1) the spatial location of the peak of the H$\\alpha$ emission coincides with the spatial location of the 8.4 GHz emission; (2) both, the peaks of H$\\alpha$ and 8.4 GHz emissions are directly proportional to each other; (3) the constant of proportionality was chosen as to have zero emission from the center. Although the assumptions seem reasonable, the structure is definitely not a proof, but it is very suggestive, of the existence of a circumnuclear structure in H$\\alpha$.} \\label{fig6} \\label{hi} \\end{figure}   The emission at 8.4 GHz, with the higher resolution, is complex and shows an unresolved nuclear source and several sources surrounding it. The circumnuclear radio sources are at all position angles (see Figs. 4 and 5), and do not align at all with the lobes or with any other optical structure in the galaxy.  This implies that the circumnuclear sources are probably located in the plane of the disk, and might be identified with a circumnuclear structure at distances of 125 pc to 325 pc.  The radio emission from the inner $4''.5$ is then most likely a mixture of synchrotron and free-free emission and the sources are most probably giant HII regions near an inner Lindblad resonance (as it is the case in other barred galaxies).  An additional hint for the existence of a circumnuclear structure is provided by the H$\\alpha$ emission. As mentioned earlier NGC 3367 shows H$\\alpha$ emission from the central region which is unresolved from ground observations (\\cite{gar96a,gar96b}) with a bright emission from a compact source and weak extended emission from within $6''$. Assuming that the peak of H$\\alpha$ coincides with the peak of the 8.4 GHz radio continuum emission, we then have subtracted the 8.4 GHz radio continuum emission from the H$\\alpha$ emission in such a way as to have zero emission from the center. The image, H$\\alpha$ - 8.4 GHz radio continuum, that is produced is shown in Figure 6; it shows an H$\\alpha$ circumnuclear structure. This H$\\alpha$ image provides {\\it no} proof of the existence of a circumnuclear structure since we have made several assumptions: (1) that the spatial position of peak of the H$\\alpha$ emission coincides with the position of the 8.4 GHz emission; (2) that the radio continuum 8.4 GHz is related to the H$\\alpha$ emission in such a way that they are directly proportional to each other through a constant; (3) that the constant of proportionality was chosen in such a way as to have zero emission from the very center; (4) that the final subtracted image represents emission from the disk of the galaxy where we think the structure lies (near a ILR). The image however is very suggestive of the existence of such a structure. If the circumnuclear structure represents regions of massive star formation one could estimate the rate of supernova using the relation  (\\cite{con92}):  \\begin{equation} \\frac {\\nu_{SN}}{yr^{-1}}  \\sim \\frac {L_N / (10^{22} W Hz^{-1}) }  {13 (\\nu /GHz)^{-\\alpha} } \\end{equation}  using the integrated flux within the innermost $3''.5$ at 1.4 GHz and assuming that all of this emission is of non-thermal origin (which we know is wrong but hope that not by much) and a spectral index $\\alpha=-0.53$ then the expected supernova rate in the circumnuclear region of NGC 3367 is $\\nu_{SN}=0.03$, which is very similar for other galaxies (\\cite{con92}). The global star formation rate using the far infrared ({\\it IRAS}) luminosity and using the relation (\\cite{con92}):  \\begin{equation} \\frac {SFR(M\\geq5 M_{\\odot})}{M_{\\odot}yr^{-1}} \\sim \\frac{L_{FIR}/L_{\\odot}}{1.1\\times10^{10}} \\end{equation}  with a L$_{FIR}=2\\times10^{10} L_{\\odot}$ (\\cite{soi89}) we get SFR$_{NGC 3367}\\sim1.8 M_{\\odot}~yr^{-1}$ a value which similar to the values found in other galaxies (\\cite{con92}).   As stated above, the gas in the circumnuclear structure could have been driven inwards either by the perturbations induced by the potential of the stellar bar, or by a possible off-center collision with a minor intruder. A reason in favor of the active compact radio nucleus is the short extension in the 8.4 GHz map that has a P.A. similar to that of the jet-like structure observed in the 1.4 GHz map that connects the central radio source with the SW lobe. If the flow were coming from a starburst wind (originating from the circumnuclear structure), this wind would be directed (in the most simple case) out of the plane towards the rotation axis of the disk, (where the density gradient is largest) and the jet - like and lobes would be projected in the SE - NW direction as mentioned above. Therefore, although an elongated structure is not observed at 8.4 GHz forming a jet, the observations suggest that the plasma flows out from the compact nucleus, possibly directed in the observed orientation as a result of an accretion disk inclined (but not perpendicular) with respect to the plane of the galaxy. One thus could infer that the jet does not interact with the circumnuclear material. The interpretation of the origin of the energy for the jets, in NGC 3367, might still controversial in the following respect, namely, the global q parameter. The q parameter is the ratio of the FIR ({\\it IRAS}) emission to the 1.4 GHz emission (\\cite{hel85}). Empirically the median $2.2 \\leq \\langle q \\rangle \\leq 3.1$ was found for spiral galaxies (\\cite{con95}), while $\\langle q \\rangle \\leq 2$ was found for galaxies powered by an AGN (\\cite{con91b}). In particular, based on the value found for q for NGC 3367 (q$=2.4$ using the 5 GHz total emission) it was concluded that the dominant energy source in NGC 3367 is stars (\\cite{con91b,con95}). If we compute q using the total emission at 1.4 GHz (119 mJy [\\cite{con98}]), we get is q=1.9 which might indicate the presence of an (mildly) AGN according to the convention (\\cite{con91b,con95}). We believe that this result involves global emissions (radio and infrared) and not necessarily the energy powering the jets and lobes observed in NGC 3367 (\\cite{con91b,con95}). This result still needs to be confronted with the current observations of the central radio sources, jet and lobes in NGC 3367.  This central radio continuum structure (compact source with circumnuclear regions) is very similar to the structure observed in some barred Sy 1 galaxies, like NGC 1097 (\\cite{hum87}) and NGC 7469 (\\cite{con91a,wil91b,mil94,mau94,gen95}).  The radius of the circumnuclear structure in NGC 1097 is about 550 pc (\\cite{hum87}), and about 450 pc in NGC 7469 (\\cite{con91c}).  These are similar to the one observed in NGC 3367 ($\\sim 300$ pc).  None of these Seyfert galaxies, however, show any indication of a large scale jet or radio continuum extended lobes (\\cite{hum87,wil91b}). In the case of normal barred galaxies, like NGC 1326 and NGC 4314, they show H$\\alpha$ circumnuclear structures with radio continuum emission characteristic from regions of star formation at similar radii (400 pc and 325 pc, respectively). None of them have radio continuum emission from the compact nucleus (\\cite{gar91a,gar91b}) nor any radio continuum extended lobes. These differences in radio continuum emission morphology from barred galaxies probably suggest also differences in disk properties, like the amount of gas in the disk, the strength of the gravitational potential including the non-axisymmetric component, and possibly different ways of transporting material to the central regions and compact nucleus.  Our main results are as follows:  1) There is a faint structure that connects the central source with the SW lobe, here identified as a low surface brightness jet.  Thus, the lobes are currently being fed by plasma from the compact radio source.  2) At the highest resolution, the 8.4 GHz emission shows an unresolved central peak with several circumnuclear sources.  The unresolved source, which is located at the center of the galaxy, is $\\sim$ 10 times stronger than individual peaks of emission in the circumnuclear region, and its deconvolved diameter is smaller than 65 pc.  3) The circumnuclear radio sources are most likely not associated with the interaction of the jet and the sourrounding medium.  4) An estimate of the spectral index from the innermost region within  $3''.5$ is about -0.53.  5) The flow of plasma from the nuclear region to the lobes is out of the plane of the galaxy but inclined with respect to the rotation axis of the disk.  6) No emission is detected from the stellar bar.  Our high resolution imaging has confirmed ideas from earlier work that NGC3367 is a currently active low power radio galaxy, continuing to be powered by a weak Liner/HII nucleus. Like other barred spirals, it shows indications of star formation both in the center and on larger scales. The jet-like structure connecting the central source with the SW lobe is out of the plane of the disk."
    },
    "0201/astro-ph0201468_arXiv.txt": {
        "abstract": "{ Systematic {\\it Rossi X-Ray Timing Explorer (RXTE)} observations of the Small Magellanic Cloud (SMC) have revealed a previously unknown transient X-ray pulsar with a pulse period of 95s. The 95s pulsar, provisionally designated XTE SMC95, was detected in three Proportional Counter Array (PCA) observations during an outburst spanning 4 weeks in March/April 1999. The pulse profile is double peaked reaching a pulse fraction of $\\approx$0.8. The X-ray spectrum is well represented by an absorbed power-law with a photon index of 1.4 and mean unabsorbed flux of $\\ga$ 8.95$\\times 10^{-11}$ ergs cm$^{-2}$ s$^{-1}$ (3 - 25 keV). The source is proposed as a Be/neutron star system on the basis of its pulsations, transient nature and characteristically hard X-ray spectrum. The 2-10 keV X-ray luminosity implied by our observations is $\\ga$ 2$\\times$10$^{37}$ergs s$^{-1}$ which is consistent with that of normal outbursts seen in Galactic systems. This discovery adds to the emerging picture of the SMC as containing an extremely dense population of transient high mass X-ray binaries. ",
        "introduction": "Observations of the SMC by {\\it RXTE, ASCA} \\& {\\it SAX} have detected an increasing number of new transient X-ray pulsars over the last few years, e.g. Haberl \\& Sasaki (\\cite{hs00}), Yokogawa et al (\\cite{Yokogawa00}). With no sign of a decline in this trend, the picture emerging from this flurry of new identifications and rediscoveries of sources seen by {\\it Einstein} and {\\it ROSAT} is that the SMC contains a large population of X-ray binary systems, of which only a small fraction are active at any one time. The SMC provides the closest available approximation to a luminosity-limited sample of X-ray pulsars due to its low extinction and small size. The depth of the SMC is small relative to its distance from us, placing all of the pulsars at effectively the same distance.  {\\it RXTE}'s proportional counter array is well suited to studying this population since its field of view, sensitivity and timing resolution allow simultaneous monitoring of all sources in a significant fraction of the SMC.  Observational data from across the spectrum has provided a coherent classification of X-ray binaries, both in the Galaxy and Magellanic clouds. The principal distinction is between high and low mass stellar counterparts, forming a distinction that is also generally coincident with that between sources exhibiting coherent X-ray pulsations, and those that do not. In high mass X-ray binary (HMXB) systems the optical emission is dominated by the star, while in the latter low mass (LMXB) case, emission from an accretion disk around the compact object dominates. The HMXBs fall into only two subclasses on the basis of their optical classification, those with an OB supergiant (SG) optical counterpart and those involving a Be emission-line star, which account for the majority of HMXBs. These two subgroups broadly coincide with two distinct modes of X-ray behaviour, persistent and transient sources.  Systems containing a supergiant are persistent sources of X-rays which are usually modulated by eclipses due to the short orbital period and correspondingly small size of the system. In very short period systems the primary overfills its Roche-lobe and the overflowing matter is forced by its high angular momentum into a stable accretion disk around the neutron star. In more widely separated systems, persistent emission arises by accretion from a roughly spherically symmetric stellar wind. Pulse-timing measurements suggest that long-term stable accretion disks do not form in these systems, although short-lived torque-reversals observed in many systems (Bildsten et al \\cite{Bildsten97}) imply the formation of transient accretion disks.  Transient X-ray sources account for the majority of HMXBs. The neutron star's (NS) optical companion is generally a Be III-V star exhibiting a strong infrared flux excess. Long term monitoring of these systems has resulted in a `standard model' in which the Be star, rotating close to its breakup speed, is surrounded by a dense equatorial disk of cool material. This disk is prone to successive phases of building and ejection on timescales of 1-3 years (Coe \\cite{Coe00a}). The neutron star, in a wide (few 10s - 100s of solar radii) eccentric orbit around the Be star, periodically intercepts the ejected material leading to the formation of an accretion disk and X-ray emission. This general model explains the regular X-ray outbursts of Be/NS systems, which are modulated at the orbital period and normally occur close to periastron passage of the neutron star. The situation is complicated by changes in the circumstellar disk which alter the density of material being encountered by the neutron star. This leads to long epochs (several orbital periods) of heightened X-ray activity often accompanied by strong torque reversals and prolonged periods of spin-up. Over the long-term, such activity leads ultimately to a dynamic equilibrium between the orbital and neutron star spin periods since the pulsar experiences negative as well as positive torques. The equilibrium is manifested observationally in the P$_{Pulse}$ vs P$_{Orbit}$ diagram (Corbet et al \\cite{Corbet99}). See also Corbet (\\cite{Corbet86}), Stella et al (\\cite{Stella86}) and Bildsten et al (\\cite{Bildsten97}) for detailed explanations.  With the addition of recent discoveries there are now 25 X-ray pulsars known in the SMC. Of these X-ray sources, the majority have been identified as belonging to Be systems rather than supergiant systems (Haberl \\& Sasaki \\cite{hs00}). Comparing these figures to the population in the Galaxy in which about 70 HMXBs are known, 31 identified with Be optical companions, the SMC emerges as containing a particularly dense population of Be X-ray binaries. Furthermore of the 14 supernova remnants (SNR) identified in the SMC none is Crab-like (Yokogawa et al \\cite{Yokogawa00}). Since Crab-like SNR are believed to form from the explosion of a single massive star it is likely that the evolutionary history of the SMC has resulted in a large ratio of massive binary stars to isolated massive stars. When the comparison is scaled by relative masses of the SMC and Galaxy it is clear that the population density of HMXBs in the SMC is higher even than the Galaxy's most populous regions. Many authors e.g. Stavely-Smith et al (\\cite{Stavelysmith97}) and Popov et al (\\cite{Popov98}) have proposed an evolutionary scenario for the SMC in which an extremely active epoch of star formation occurred relatively recently (within 5 Myr), leading to the large proportion of young, massive stars. In support of this theory are 4 key observations: (1) Large number of supernova remnants approximately 5 Myr old; (2) The large number of HMXBs; (3) Apparent absence of low mass systems which would represent an older stellar population. (These still remain elusive despite the current generation of satellites' ability to detect them if they are similar to galactic LMXBs); (4) The ratio of Be/B type stars in the SMC is 0.39 (Maeder et al \\cite{mgm99}) compared to 0.16 in the Galaxy, implying a younger population. ",
        "conclusions": "Most galactic Be/X-ray binaries increase in luminosity by a factor of $\\ga$100 during normal outburst (Negueruela \\cite{Negueruela98}) reaching luminosities around 10$^{37}$erg s$^{-1}$. Taking the distance to the SMC as 63 kpc (Groenewegen et al \\cite{Groenewegen00}), the fluxes given in Table~\\ref{spec} imply a 2-10 keV luminosity for XTE SMC95 of at least 2$\\times$10$^{37}$ erg s$^{-1}$. Thus during this outburst the X-ray luminosity of XTE SMC95 was comparable to that of Galactic Be/NS systems. The photon indices obtained from our spectral fits are slightly steeper than those normally seen in Galactic Be/NS HMXBs, values $\\la$ 1.0 being typical in the 10-20 keV range (White et al \\cite{White83}). Softer (steeper) spectra have been associated with less luminous sources based on studies in the Galaxy (White et al \\cite{White83}). Comparing our results (Table~\\ref{spec}) to the recent outburst of SMC X-2 which was both more luminous and had a harder spectrum: luminosity $\\approx$10$^{38}$ergs s$^{-1}$, spectral index 0.7-1.0 (Corbet et al \\cite{Corbet01}) this trend looks likely to hold in the SMC.  XTE SMC95 exhibited varying structure in its pulse profile. Pulse profiles obtained in two separate energy bands 3-10 keV and 10-20 keV shown in Figure~\\ref{profiles95} indicate the pulse-shape changed between the two observations (3 \\& 4), with the second pulse peak weakening relative to the first, accompanied by a decrease in pulsed flux. The significant shift in relative strength of the first and second power spectrum harmonics between the two observations (see Figures~\\ref{powspec1} \\&~\\ref{powspec2}) confirms that the profile change is real and not an artifact of an incorrect profile-folding period.  A small decrease in the pulse period was detected between the two observations of XTE SMC95 one week apart, implying a pulse period derivative (8.6 $\\pm$ 5.5) $\\times$10$^{-7}$ s s$^{-1}$. This change is attributed to an unknown combination of spin up and orbital motion depending on the inclination of the system.  For Be/NS systems with P$_{pulse}$ around 95s, the Corbet diagram (Corbet \\cite{Corbet86}) suggests likely orbital periods in the 50-100d range. XTE SMC95 was detected for three weeks, only a small fraction of its likely orbital period. Whether this outburst was related to orbital motion or long-term changes (disk-loss events) in a circumstellar disk may be determined by continued monitoring.  A search of the ROSAT source catalogue (Haberl et al \\cite{hfp00}) reveals 4 sources within the 3$\\sigma$ confidence region (See Figure~\\ref{errorbox}). Of these sources, RX J0053.8-7252 is classified as a candidate X-ray binary and RX J0052.6-7247 a background AGN the other sources are a foreground star and a Supersoft X-ray source. A similar search for known Be stars was made in the same region. Searching the AzV (Azzopardi \\& Vigneau \\cite{azv82}) \\& LIN (Lindsay \\cite{lin}) catalogues for spectral types: O8-B2, I-V, gave 7 stars within the 3$\\sigma$ contour. This range of spectral types was chosen to match the known distribution for optical counterparts to HMXBs (Negueruela \\cite{Negueruela98}). The AzV catalogue is 80\\% complete to limiting magnitude B=14.3, considering the distance to the SMC the sample is strongly biased toward early-type giants. Thus main sequence Be stars (Luminosity class IV-V) are likely to be under represented in the sample. The catalogue of H$\\alpha$ emission-line objects of Meyssonnier \\& Azzopardi (\\cite{ma93}) lists 69 stars within the 3$\\sigma$ contour. From the argument given above, the majority of these are likely to be O/Be stars."
    },
    "0201/astro-ph0201142_arXiv.txt": {
        "abstract": "This paper describes a framework for studying galaxy morphology, particularly bar strength, in a quantitative manner, and presents applications of this approach that reveal observational evidence for secular evolution in bar morphology.  The distribution of bar strength in galaxies is quite strongly bimodal, suggesting that barred and unbarred systems are distinct entities, and that any evolution between these two states must occur on a relatively rapid timescale.  Bars' strengths appear to be correlated with their pattern speeds, implying that these structures weaken as they start to slow, and disappear entirely before the bars have slowed significantly.  There is also tantalizing evidence that bars are rare beyond a redshift of $z \\sim 0.7$, indicating that galaxies have only recently evolved to a point where bars can readily form. ",
        "introduction": "I had originally intended to give an overview of the observational evidence relating to the mechanisms by which bars evolve over time. However, this overview has been expertly provided by a number of other authors in these proceedings.  Bureau (2002), for example, gives a thorough treatment of the buckling of bars to form boxy bulges, and how this may lead to their dissolution.  Similarly, Das et al.\\ (2002) discuss the role that central mass concentrations, such as massive black holes, may play in weakening and ultimately destroying bars. Debattista (2002) describes the way that bars might be expected to slow down through dynamical friction, while Gerssen (2002) presents the observational evidence that not much slowing has occurred.  This apparent contradiction suggests that the mass distribution in barred galaxies cannot be very centrally concentrated, or that some other mechanism destroys the bars before they have had the opportunity to slow down.  Finally, Athanassoula (2002), through her sophisticated N-body simulations, discusses much of the theory behind these observations.  This wide sweep of articles leaves me very little to say by way of background review.  I will therefore limit this paper to discussing how these observations, and the accompanying theory, might be brought together in a common quantitative framework, and to presenting a couple of preliminary applications of this procedure.  Section~\\ref{sec:quant} discusses the quantification of bar strength, and how it can be combined with other quantitative measures of morphology to define a parameter space directly analogous to that used in the previous qualitative classification of Hubble's tuning fork. Section~\\ref{sec:Omegap} describes an application of this parameterization in studying the evolution of bar pattern speeds, and Section~\\ref{sec:HDF} looks at direct evidence for bar evolution in the Hubble Deep Fields.  Finally, Section~\\ref{sec:conc} speculates on the next steps in this field. ",
        "conclusions": "\\label{sec:conc}  With the advent of largescale CCD surveys of galaxies, both nearby and in the more distant Universe, the study of galaxy morphology has matured from the qualitative to the quantitative.  Robust techniques have been developed for automatically measuring properties of galaxies such as their bulge-to-disk ratios and bar strengths.  Perhaps unsurprisingly, these quantitative measures correspond quite closely to their qualitative predecessors, to the extent that one can reproduce Hubble's tuning fork classification.  The bifurcation between barred and unbarred galaxies has significant implications for investigations of secular evolution in bars, since it implies that any evolution between these two end states must occur on a timescale that is short compared to the galaxies' cosmological lifetimes.  Further evidence for evolution in bar strength on timescales that are less than the cosmological have come from studies of faint galaxies in the HDFs, which indicate that bar formation is a relatively recent phenomenon.  An important next step will be to use this framework to compare the observations with the predictions of N-body simulations.  It is now possible to calculate morphological parameters to follow the properties of N-body simulations, such as those described by Athanassoula (2002), as the systems undergo secular evolution. Further, cosmological simulations like those presented by Navarro (2002) are rapidly advancing to the point where they can reliably follow the formation and evolution of individual disk galaxies, so we can also follow the variations in morphological parameters in a cosmological framework.  Through such analysis, galaxy observations such as those presented here can be directly tied to N-body simulations to test our understanding of the formation and evolution of these systems, just as the observed sequences of stars in the color-magnitude diagram can be combined with stellar evolutionary tracks in order to understand how stars form and evolve.  One other area where there is still much room for analysis lies in the morphological description of bars.  As described in Section~\\ref{sec:quant}, it seems unlikely that the essence of a bar can be captured in a single parameter, but it is by no means clear what combination of parameters does provide the optimal description, nor what the physical differences are that give rise to the range of parameters observed.  A fuller description of morphological features such as bars will enable us to bring more sophistication to our exploration of the mechanisms by which these features evolve, and hopefully answer many of the outstanding questions."
    },
    "0201/astro-ph0201374_arXiv.txt": {
        "abstract": "The 3rd Interplanetary Network (IPN) has been operating since April 2001 with two distant spacecraft, Ulysses and Mars Odyssey, and numerous near-Earth spacecraft, such as BeppoSAX, Wind, and HETE-II. Mars Odyssey is presently in orbit about Mars, and the network has detected approximately 30 cosmic, SGR, and solar bursts.  We discuss the results obtained to date and use them to predict the future performance of the network. ",
        "introduction": "The 3rd IPN began with the launch of Ulysses in November 1990.  Ulysses is in a heliocentric orbit roughly perpendicular to the ecliptic, with perihelion  of about 1.5 AU and an aphelion of about 5 AU.  Until 1992, the network had Ulysses and Pioneer Venus Orbiter (PVO) as its distant points, and utilized many near-Earth spacecraft such as the Compton Gamma-Ray Observatory (CGRO) as its third point.  PVO entered the atmosphere of Venus in 1992, and it was to be replaced in the IPN by NASA's Mars Observer, which was lost during insertion into Martian orbit.  Finally, in 1999, the network was completed by the X- and Gamma-Ray Spectrometer experiment aboard the Near Earth Asteroid Rendezvous (NEAR) spacecraft.  In this configuration, the IPN operated quite successfully until NEAR landed on the asteroid Eros in February 2001.  Figure 1 summarizes some of the results. The Mars Odyssey mission, launched in April of that year, contains two experiments which have been modified for gamma-ray burst detection, the Gamma-Ray Spectrometer (GRS) and the High Energy Neutron Detector (HEND).  As the GRS will not be turned on permanently until the spacecraft completes its aerobraking maneuver, we will discuss the results obtained with HEND.  \\begin{figure} \\psfig{file=figure1.eps,width=10cm} \\caption{Fifty-seven GRBs detected by the IPN with NEAR.  Each burst is characterized by the size of the error box and the delay to obtain it.  The stars represent GRBs for which counterparts were identified; six of them have had their redshifts measured.  The filled circles represent bursts which were followed up in the radio and/ or optical range, but for which no counterparts were identified.  The hollow circles represent those events which were not followed up.  In slightly over a year of operation, the database on GRB counterparts increased by 50\\% due to the IPN results alone.} \\label{fig1} \\end{figure} ",
        "conclusions": ""
    },
    "0201/astro-ph0201004_arXiv.txt": {
        "abstract": "We consider the capabilities for detecting low order CO emission lines from high-redshift ($z$) galaxies using the next generation of radio telescopes operating at 22 and 43\\,GHz. Low order CO emission studies provide critical insight into the nature of high redshift galaxies, including: (i) determining molecular gas masses, (ii) study of large scale structure through 3-dimensional redshift surveys over cosmologically relevant volumes, (iii) imaging gas kinematics on kpc-scales,  and, in conjunction with observations of higher order transitions using future millimeter telescopes, (iv) constraining the excitation conditions of the gas. Particular attention is paid to the impact on such studies of the high frequency limit for future centimeter telescopes.  We employ models for the evolution of dusty star forming galaxies based on source counts at (sub)millimeter (mm) wavelengths, and on the observed mm through infrared (IR) backgrounds, to predict the expected detection rate of low-order CO(2-1) and CO(1-0) line emitting galaxies for optimal centimeter(cm)-wave surveys using future radio telescopes, such as the Square Kilometer Array (SKA) and the Expanded Very Large Array (EVLA). We then compare these results to surveys that can be done with the next-generation mm-wave telescope, the Atacama Large Millimeter Array (ALMA). Operating at 22\\,GHz the SKA will be competitive with the ALMA in terms of the detection rate of lines from high-$z$ galaxies, and will be potentially superior by an order of magnitude if extended to 43\\,GHz. Perhaps more importantly, cm-wave telescopes are sensitive to lower excitation gas in higher redshift galaxies, and so provide a complementary view of conditions in high redshift galaxies to mm-wave surveys. We have also included in our models emission from HCN. The number of HCN(1-0) detections will be about 5$\\%$ of the CO detections in the (CO-optimized) 22 GHz surveys, and about 1.5$\\%$ for 43 GHz surveys. In order not to over-resolve the sources, brightness temperature limitations require that a future large area cm telescopes have much of its collecting area on baselines shorter than 10 km. ",
        "introduction": "The models we employ for predicting CO source counts are updated from Paper I. These models employ an analytic description of pure luminosity evolution of the low-$z$ {\\it IRAS} 60-$\\mu$m luminosity function (Saunders et al.\\ 1990). The evolution function has the form, \\begin{equation} g(z) = (1+z)^{3/2} {\\rm sech}^2 [ b {\\rm ln}(1+z) - c ] \\, {\\rm cosh}^2 c. \\end{equation} At very low and very high redshifts the function can be approximated by $g(z) \\propto (1+z)^{\\gamma}$, with $\\gamma \\simeq {3/2} + 2b\\,{\\rm tanh}\\,c$ and $\\gamma = {3/2} - 2b$ respectively. With the current cosmology, and taking into account all available far-IR and (sub)-mm background, count and redshift distribution data, the values $b = 2.2 \\pm 0.1$ and $c = 1.84 \\pm 0.15$ are required, if a typical dust temperature of 37\\,K is assumed. The fitting procedures and details of the information used are explained in Blain et al.\\ (1999). A plot of this function is shown in Fig.\\,1 of Blain (2002). The evolution peaks at $z \\simeq 1.7$, at which the bolometric luminosity density of infrared-luminous galaxies is about 40 times greater than at $z=0$.  The strength of the emission from CO lines is calculated as described in Paper I. Excitation conditions, and thus line ratios, are derived from a  standard large velocity gradient model (Frayer \\& Brown 1997) with a kinetic temperature of about 50\\,K and a density of 10$^4$\\,cm$^{-3}$. This model provides a reasonable description of the CO emission from two well-studied FIR-luminous galaxies at $z \\sim 2.5$, and their properties are used to normalize the CO emission line strengths to the evolving bolometric luminosity function of galaxies described above. ",
        "conclusions": "From Table\\,1 and Fig.\\,2, it can been seen that future large area cm-wave telescopes  operating at 22 GHz  will be competitive with future mm-wave telescopes in terms of discovering high-$z$ star-forming galaxies through their molecular line emission. Increasing the high frequency limit to 43-GHz allows for optimal surveys which are 20 times faster than the ALMA in terms of discovering high-$z$ galaxies.  Detection rates are a simplistic metric for the ALMA and the SKA, and it is important to emphasize their complementarity. ALMA surveys will be dominated by higher-order CO lines from intermediate redshift ($z \\sim 1 - 2$),   intermediate luminosity (L$_{\\rm FIR} \\sim ~ \\rm few\\times10^{11}$ L$_\\odot$) objects. Surveys with the SKA at 22 GHz will be dominated by low-order transitions from higher luminosity (L$_{\\rm FIR} \\sim 10^{12}$ L$_\\odot$) galaxies at higher  redshifts ($z \\sim 4$).  For the EVLA at 43 GHz the combined small FoV and low sensitivity make it a much slower survey instrument that either the ALMA and SKA. However, the EVLA has adequate sensitivity to resolve and study at sub-arcsec angular resolution the low-order CO emission from individual FIR-luminous, high-$z$ objects selected from wide-field surveys at other wavelengths, such as surveys using (sub)mm bolometer arrays on large single-dish telescopes. And perhaps most importantly, the large fractional bandwidth of the EVLA will allow for redshift determinations via molecular line searches using pointed observations of individual objects. The EVLA will thus provide the first look into the nature of low-order CO emission from high-$z$ galaxies. The importance of such capabilities has already been demonstrated with the current VLA in a few extreme cases, although the limited bandwidth effectively precludes  proper spectroscopy (Papadopoulos et al.\\ 2001; Carilli et al.\\ 1999, 2002).  The CO excitation conditions assumed in the models of section 2 could lead to pessimistic predictions for the low-order transitions because the velocity integrated line flux density increases roughly with the square of the frequency (ie. constant brightness temperature), at least to CO(4-3) (Paper I). While this appears to be roughly appropriate for infra-red selected galaxy samples, our models do not include the possibility of a population of molecular gas-rich, high redshift galaxies with lower excitation conditions. For example, the CO excitation conditions for the Milky Way disk inside the solar radius (excluding the Galactic center) imply roughly equal velocity integrated flux density for CO(1-0) and CO(4-3) (Fixsen, Bennett, \\& Mather 1999). If such a population of galaxies exists, then the predictions from our models can be considered lower limits to the cm-wave source counts for molecular line surveys. Recent observations with the VLA provide evidence that such a population may indeed exist (Papadopoulos et al.\\ 2001).  Limitations to the intrinsic brightness temperature of the thermal line emission from high redshift galaxies require that much of the collecting area of future large area radio telescopes be concentrated on baselines $\\le 10$ km.  On the other hand, having baselines out to 10 km provides the important capability of imaging the emission on scales relevant to galaxies ($0.14'' \\simeq 1$ kpc), and for resolving the multiple images of the order of 1\\% of line-emitting galaxies expected to be gravitationally lensed by foreground galaxies.  Lastly, we have found that HCN(1-0) emission will not be a major source of confusion to optimal cm-wave CO line searches, comprising about 5$\\%$ of the total number of detected galaxies at 22 GHz and 1.5$\\%$ at 43 GHz. Of course, once identified as such, HCN is interesting in it's own regard as a better tracer than CO of the star forming clouds in active star forming galaxies (Solomon 2001).  Like studies of star formation in our own galaxy, it has become clear that a complete census of the star-formation history of the universe requires an understanding of the contribution from galaxies that are obscured by dust at optical wavelengths. The next generation mm- and cm-wave telescopes will provide unique, and complementary,  capabilities for studying the  thermal and non-thermal line and continuum emission from such systems at sub-arcsecond spatial resolution.  \\vskip 0.2truein  The National Radio Astronomy Observatory (NRAO) is operated by Associated  Universities, Inc. under a cooperative agreement with the National Science Foundation. We thank Dave Frayer for further discussions  about the line models implemented in Paper I."
    },
    "0201/astro-ph0201232_arXiv.txt": {
        "abstract": "We have carried out a high-dispersion (R$\\sim$30,000) echelle spectroscopic survey of 16 Type II supernovae (SNe) to search for narrow emission lines from circumstellar nebulae ejected by their massive progenitors.  Circumstellar nebulae, if detected, provide invaluable opportunities to probe SN progenitors.  Of the 16 SNe observed, SN ejecta are clearly detected in 4 SNe and possibly in another 2 SNe, interstellar gas is detected in 12 SNe, and circumstellar material is detected only in SN\\,1978K and SN\\,1998S. In the case of SN\\,1978K we are able to place an upper limit of $\\sim$2.2 pc for the size of the circumstellar ejecta nebula and note that this is more consistent with the typical sizes observed for ejecta nebulae around luminous blue variables rather than Wolf-Rayet stars.  In the case of SN\\,1998S, our observations of the narrow lines $\\sim$1 year after the SN explosion show variations compared to early epochs. The nebular lines we observe from SN\\,1998S originate from either the low--density, outer region of a circumstellar nebula or have become dominated by an interstellar component. ",
        "introduction": "Supernovae (SNe) of Types Ib, Ic, and II are believed to have massive progenitors, because they have been found frequently in or near spiral arms and \\ion{H}{2} regions but not in elliptical galaxies \\citep*{VDHF96}. Few SNe (e.g., SN\\,1987A) have recognized massive progenitors; consequently, little is known about the stellar evolution immediately before the SN explosion.  Evolved massive stars are known to undergo copious mass loss, forming circumstellar nebulae. The rings around SN~1987A are an example of such a nebula \\citep{Buetal95}.  The chemical composition and kinematics of the rings have provided essential constraints that lead to the hypothesis that the B3I progenitor Sk\\,$-$69$^\\circ$202 was a binary \\citep{Po92}.  These circumstellar nebulae can also be detected from the presence of narrow H$\\alpha$ emission lines (FWHM $\\leq$ 200 km s$^{-1}$) in spectra of Type IIn SNe \\citep{Sc90,filip91,filip97}.  The circumstellar nebulae of distant SNe cannot be resolved spatially, but their expansion, physical conditions, and chemical enrichment can be investigated by high-dispersion (R $\\geq$ \\ 30,000) spectra, as recently demonstrated for SN\\,1997ab \\citep{sal98} and SN\\,1978K \\citep{Chetal99}.  Most available spectroscopic observations of SNe have been made with low or intermediate spectral dispersion to study the broad spectral lines of SN ejecta.  These spectra are useful for revealing the presence of narrow nebular lines, but are not adequate for analysis of nebular kinematics and physical conditions.  Therefore, we have undertaken a high-dispersion spectroscopic survey of 16 Type II SNe.  The results are reported in this paper. ",
        "conclusions": "\\label{discuss}  The very last evolutionary stage of a massive star before SN explosion is not well known. It is impossible to study SN progenitors after they have exploded, but it is possible to study the CSM shed by the progenitors and photoionized by UV flashes from the SNe.   The physical properties of these circumstellar nebulae may provide invaluable information about the doomed massive stars.  In the last hundred years, SN\\,1987A has been the only SN observed in the Local Group.  All the other SNe are in distant galaxies, where circumstellar nebulae can only be detected in high-dispersion spectra of SNe as narrow nebular lines with high [\\ion{N}{2}]/H$\\alpha$ ratios.  The size of an unresolved circumstellar nebula can be determined if the SN is spectroscopically monitored until the SN ejecta impact the circumstellar nebula.  For example, if SN\\,1987A were at a large distance, the size of its inner ring nebula could still be determined from the 10--11 years time lapse from the SN explosion to the emergence of a broader nebular component, and the 15,000 km~s$^{-1}$ expansion velocity of the SN ejecta inferred from its broad Ly$\\alpha$ emission \\citep{Mietal00}.  There are apparently two types of circumstellar nebulae produced by SN progenitors.   The first type of circumstellar nebulae are swept-up by SN ejecta within a year or two, as demonstrated dramatically in time-sequenced spectra for SN\\,1988Z \\citep{stat91}.  In the case of SN\\,1997eg, previous observations have shown a narrow P-Cygni component to the SN spectrum \\citep{sal02} which has disappeared at the time of our observations, one year later. These nebulae are small ($\\sim$0.01 pc), and, for the case of SN\\,1997eg, very dense ($\\ge 10^7$ cm$^{-3}$). Such ionized nebulae have not been observed around any known evolved massive stars.  This circumstellar material may have been ejected immediately before the SN explosion \\citep{Chu01} or during a red supergiant phase \\citep{fassia01}.  The second type of circumstellar nebulae are longer lived, indicating a larger size.  These nebulae might be the counterparts of circumstellar nebulae observed around Wolf-Rayet (WR) stars or LBVs \\citep{CWG99}.  For a nebular radius of 2 pc and a SN ejecta expansion velocity of 10,000 \\kms, the impact of SN ejecta on the circumstellar nebula is expected at $\\sim$200 yr after the SN explosion, and the impact will produce X-ray-luminous SNRs such as the one observed in NGC\\,6946 \\citep{dunne00} and the SNR 0540$-$69.3 in the Large Magellanic Cloud \\citep{CMB98}.  Among the SNe we surveyed, we detect CSM around SN\\,1978K and SN\\,1998S.  In the case of SN\\,1978K, the circumstellar nebula was not yet hit by the SN ejecta 21 years after the SN explosion. This must belong to the second type of circumstellar nebula described above.  In the case of SN\\,1998S, the rapid temporal evolution of the spectral lines from the CSM over the first year indicate that this material is interacting with the SN ejecta and must belong to the first type.   Our new observations of SN\\,1978K confirm the \\ha\\ and \\nii\\ line emission and expansion velocity of the CSM previously reported by \\citet{Chetal99}.  As this circumstellar nebula is in the second class mentioned above and NGC\\,1313 is relatively nearby, the nebula might be large enough to be resolved by the {\\it HST}. Therefore, we obtained the {\\it HST} archive WFPC2 images of SN\\,1978K taken with F656N and F606W filters on 1998 September 23 as part of program GO-6713 (PI: W. Sparks). In these observations the SN is centered in the PC and a point-like source is detected (see Figure~\\ref{fig_psf}).  Also shown in Figure~\\ref{fig_psf} are the expected point-spread-functions for position of the SN in the PC for observations with the F656N and F606W filters.\\footnote{The theoretical PSFs were generated using Tiny Tim V5.0 written by John Krist and Richard Hook which is available from http://www.stsci.edu/software/tinytim.} These images clearly place an upper limit of 0\\farcs 1 on the size of the circumstellar nebula of SN\\,1978K. At the distance of the host galaxy NGC\\,1313, 4.5 Mpc \\citep{devauc63}, this upper limit on the size corresponds to 2.2 pc. This small size is more consistent with the sizes of LBV nebulae than WR nebulae, as also concluded by \\citet{Chetal99}.  In the case of SN\\,1998S, our observations cannot unambiguously determine whether the emission seen is circumstellar or interstellar.  In either case, the nebular emission observed is no longer consistent with gas denser than the critical density needed for the auroral \\nii\\ $\\lambda$5575 line to dominate the nebular lines. Therefore, the density must be less than $\\sim$10$^{5}$ cm$^{-3}$ and the emission seen is either from a lower density outer region of a circumstellar nebula or from interstellar material.  If the narrow lines we detected are indeed circumstellar in origin, then the blueward velocity offset of the \\nii\\ line relative to the \\ha\\ and \\oiii\\ lines may be caused by light travel time effects similar to those used to explain the early evolution of the \\oiii\\ line profile by \\citet{fassia01}.  Our high-dispersion (R$\\sim$30,000) spectroscopic observations of SNe have demonstrated the possibility of detecting circumstellar nebulae ejected by SN progenitors and deriving information on the progenitors from the nebular properties. Our observations have also demonstrated the difficulty of such work because of the faintness of the SNe and their associated circumstellar nebulae.  In order to identify emission from such nebulae, high-dispersion spectroscopic observations of SNe with large telescopes such as Gemini and Keck are needed.  Only after a large number of such circumstellar nebulae have been detected and studied can we hope to better understand SN progenitors and their circumstellar environment."
    },
    "0201/hep-ph0201194_arXiv.txt": {
        "abstract": "Direct detection experiments for neutralino dark matter in the Milky Way are examined within the framework of SUGRA models with R-parity invariance and grand unification at the GUT scale, $M_G$. Models of this type apply to a large number of phenomena, and all existing bounds on the SUSY parameter space due to current experimental constraints are included. For models with universal soft breaking at $M_G$ (mSUGRA), the Higgs mass and $b\\rightarrow s\\gamma$ constraints imply that the gaugino mass, $m_{1/2}$, obeys $m_{1/2} >$(300-400)GeV putting most of the parameter space in the co-annihilation domain where there is a relatively narrow band in the $m_0 - m_{1/2}$ plane. For $\\mu > 0$ we find that the neutralino -proton cross section $\\stackrel{>}{\\sim} 10^{-10}$ pb for $ m_{1/2} < 1$ TeV, making almost all of this parameter space accessible to future planned detectors.  For $\\mu < 0$, however, there will be large regions of parameter space with cross sections $< 10^{-12}$ pb, and hence unaccessible experimentally. If, however, the muon magnetic moment anomaly is confirmed, then $\\mu >0$ and $m_{1/2}\\stackrel{<}{\\sim} 800$ GeV. Models with non-universal soft breaking in the third generation and Higgs sector can allow for new effects arising from additional early universe annihilation  through the $Z$-channel pole. Here cross sections that will be accessible in the near future to the next generation of detectors can arise, and can even rise to the large values implied by the DAMA data. Thus dark matter detectors have the possibility of studying the  the post-GUT physics that control the patterns of soft breaking. ",
        "introduction": "The recent BOOMERanG, Maxima and DASI data has allowed a relatively precise determination of the mean amount of dark matter in the universe, and these results are consistent with other astronomical observations. Within the Milky Way itself, the amount of dark matter is estimated to be \\begin{equation}\\rho_{DM}\\stackrel{\\sim}{=}(0.3 - 0.5){\\rm GeV/cm^3}\\end{equation} Supersymmetry with R-parity invariance possesses a natural candidate for cold dark matter (CDM), the lightest neutralino, $\\tilde\\chi^0_1$, and SUGRA models predict a relic density consistent with the astronomical observations of dark matter. Several methods for detecting the Milky Way neutralinos exist:  (1) Annihilation of $\\tilde\\chi^0_1$ in the halo of the Galaxy leading to anti-proton or positron signals. There have been several interesting analyses of these possibilities \\cite{kane,gondolo}, but there are still uncertainties as to astronomical backgrounds.  (2) Annihilation of the $ \\tilde\\chi^0_1$ in the center of the Sun or Earth leading to neutrinos and detection of the energetic $\\nu_\\mu$ by neutrino telescopes (AMANDA, Ice Cube, ANTARES). Recent analyses \\cite{barger,bottino} indicate that these detectors can be sensitive to such signals, but for the Minimal Supersymmetric Standard Model (MSSM)  one requires $m_{\\tilde\\chi^0_1} > 200$ GeV (i.e. $m_{1/2} > 500$ GeV) and $ \\tan\\beta >10$, and for SUGRA models one is restricted to $\\tan\\beta > 35$ \\cite{barger}.  (3) Direct detection by scattering of incident $ \\tilde\\chi^0_1$ on nuclear targets of terrestrial detectors. Current detectors are sensitive to such events for $ \\tilde\\chi^0_1-p$ cross sections in the range \\begin{equation} \\sigma_{\\tilde\\chi^0_1-p} \\stackrel{>}{\\sim} 1\\times 10^{-6} {\\rm pb} \\end{equation} with a possible improvement by a factor of 10 - 100 in the near future. Future detectors (GENIUS, Cryoarray, ZEPLIN IV) may be sensitive down to $(10^{-9} - 10^{-10})$ pb and we will see that this would be sufficient to cover the parameter space of most SUGRA models.  In the following we will consider SUGRA models with R-parity invariance based on grand unification at the GUT scale $M_G\\stackrel{\\sim}{=} 2\\times10^{16}$ GeV. In particular, we will consider two classes of models: Minimal supergravity models (mSUGRA \\cite{sugra1,sugra2}) with universal soft breaking masses at $M_G$, and non-universal models with non universal soft breaking at $M_G$ for the Higgs bosons and the third generation of squarks and sleptons .  Here the gaugino masses ($m_{1/2}$)  and the cubic soft breaking masses ($A_0$) at $M_G$ are assumed universal.  SUGRA models apply to a wide range of phenomena, and data from different experiments interact with each other to greatly sharpen the predictions. We list here the important experimental constraints:  Higgs mass: $m_h > $114 GeV \\cite{higgs}. The theoretical calculation of $m_h$ still has an an error of $\\sim3$ GeV, and so we will (conservatively) interpret this bound to mean $m_h(\\rm theory) > 111$ GeV.  $b\\rightarrow s\\gamma$ branching ratio. We take a $2\\sigma$ range around the central CLEO value \\cite{bsgamma}: \\begin{equation} 1.8 \\times 10^{-4} \\leq B(B \\rightarrow X_s\\gamma) \\leq 4.5 \\times 10^{-4}\\end{equation}  $\\tilde\\chi^0_1$ relic density: We assume here  \\begin{equation} 0.02\\leq\\Omega_{\\rm DM} h^2\\leq0.25\\end{equation}  \\noindent The lower bound takes into account of the possibility that there is more than one species of DM. However, results are insensitive to raising it to 0.05 or 0.10. \\begin{figure}[htb] \\centerline{ \\DESepsf(edmt16.epsf width 10 cm) } \\caption {\\label{fig1} Corridors in the $m_0 - m_{1/2}$ plane allowed by the relic density constraints for (bottom to top) $\\tan\\beta = 10$, 30, 40, $A_0 = 0$ and $\\mu > 0$. The lower bound on $ m_{1/2}$ is due to the $m_h$ lower bound for $\\tan\\beta$ = 10, due to the $b\\rightarrow s \\gamma$ bound for $\\tan\\beta =$ 40, while both these contribute equally for $\\tan\\beta = 30$. The short lines cutting the channels represent upper bound from the $g_\\mu - 2$ experiment. [17]} \\end{figure} \\begin{figure}[htb] \\centerline{ \\DESepsf(adhs3.epsf  width 10 cm) } \\caption {\\label{fig2} Corridors in the $m_0 - m_{1/2}$ plane allowed by the relic density constraint for $\\tan\\beta = 40,\\, \\mu > 0$ and (bottom to top) $A_0 = 0,\\, -2 m_{1/2}$, $4m_{1/2}$. the curves terminate at the lower end due to the $b \\rightarrow s\\gamma$ constraint except for$ A_0 = 4 m_{1/2}$ which terminates due to the $m_h$ constraint. The short lines cutting the corridors represent the upper bound on $m_{1/2}$ due to the $g_\\mu -2$ experiment. [17]} \\end{figure}  Muon $a_\\mu = (g_\\mu -2)/2$ anomaly. The Brookhaven E821 experiment \\cite{BNL} reported a 2.6$\\sigma$ deviation from the Standard Model value in their measurement of the muon magnetic moment. Recently a sign error in the theoretical calculation \\cite{knecht,kinoshita} has reduced this to a 1.6$\\sigma$ anomaly, though recent measurements \\cite{Ajinenko} used to calculate the hadronic contribution may have raised the deviation. Since there is a great deal of more data currently being analyzed (with results due this spring) that will reduce the errors by a factor of $\\sim$2.5, we will assume here that there is a deviation in $a_\\mu$ due to SUGRA of amount \\begin{equation} 11 \\times 10^{-10}\\leq a_\\mu^{\\rm SUGRA}\\leq 75\\times 10^{-10} \\end{equation}  \\noindent We will, however, state our results with and without including this anomaly.  To illustrate how the different experimental constraints affect the SUSY parameter space, we consider the mSUGRA example:  (1) The $m_h$ and $b\\rightarrow s\\gamma$ constraints put a lower bound on $m_{1/2}$: \\begin{equation} m_{1/2} \\stackrel {>}{\\sim}(300 - 400){\\rm GeV} \\end{equation} which means $m_{\\tilde\\chi^0_1} \\stackrel {>}{\\sim}(120 - 160)$ GeV (since $m_{\\tilde\\chi^0_1} \\stackrel{\\sim}{=} 0.4 m_{1/2}$). (2)  Eq.(6) now  means that most of the parameter space is in the $\\tilde\\tau_1 -\\tilde\\chi^0_1$ co-annihilation domain in the relic density calculation. Then $m_0$ (the squark and slepton soft breaking mass) is approximately determined by $m_{1/2}$ as can be seen in Figs. 1 and 2. (3) If we include the $a_{\\mu}$ anomaly, since $a_{\\mu}^{\\rm SUGRA}$ is a decreasing function of $m_{1/2}$  and $m_0$ , the lower bound of Eq.(5) produces an upper bound on $m_{1/2}$ and the positive sign of $a_\\mu$ implies that the $\\mu$ parameter is positive. In addition one gets a lower bound on tanbeta of $\\tan\\beta > 5$. \\begin{figure}[htb] \\centerline{ \\DESepsf(aadcoan61020.epsf width 10 cm) } \\caption {\\label{fig3} $\\sigma_{\\tilde{\\chi}_{1}^{0}-p}$ for mSUGRA for $\\mu < 0$, $A_0 = 1500$ GeV, for $\\tan\\beta = 6$ (short dash), $\\tan\\beta = 8$ (dotted), $\\tan\\beta = 10$ (solid), $\\tan\\beta = 20$ (dot-dash), $\\tan\\beta=25$ (dashed). Note that the $\\tan\\beta = 6$ curve terminates at low $m_{1/2}$ due to the Higgs mass constraint, and the other curves terminate at low $m_{1/2}$ due to the $b \\rightarrow s\\gamma$ constraint [18].} \\end{figure} \\begin{figure}[htb] \\centerline{ \\DESepsf(adhs2.epsf  width 10 cm) } \\caption {\\label{fig4} $\\sigma_{\\tilde{\\chi}_1^0-p}$ as a function of the neutralino mass $m_{\\tilde{\\chi}_1^0}$ for $\\tan\\beta = 40$, $\\mu > 0$ for $A_0 = -2 m_{1/2}, 4 m_{1/2}, 0$ from bottom to top. The curves terminate at small $m_{\\tilde{\\chi}_1^0}$  due to the $b \\rightarrow s\\gamma$ constraint for $A_0 = 0$ and $- 2 m_{1/2}$ and due to the Higgs mass bound ($m_h > 114$ GeV) for $A_0 = 4 m_{1/2}$. The curves terminate at large $m_{\\tilde{\\chi}_1^0}$ due to the lower bound on $a_{\\mu}$ of Eq. (5)[17].} \\end{figure} Thus the parameter space has begun to be strongly constrained, allowing for more precise predictions. In order to carry out detailed calculations, however, it is necessary to include a number of analyses to obtain accurate results. We list some of these here:  Two loop gauge and one loop Yukawa renormalization group equations (RGE) are used in going from $M_G$ to the electroweak weak scale $M_{\\rm EW}$, and QCD RGE are used below $M_{\\rm EW}$ for the light quark contributions. Two loop and pole mass corrections are included in the calculation of $m_h$. One loop corrections to $m_b$ and $m_\\tau$ \\cite{rattazi,carena} are included which are important at large $\\tan\\beta$. Large $\\tan\\beta$  NLO SUSY corrections to $b \\rightarrow s \\gamma$ \\cite{degrassi,carena2} are included. In calculating the relic density, all stau-neutralino co- annihilation channels are included, and this calculation is done in a fashion valid for both small and large $\\tan\\beta$.  We do not include Yukawa unification or proton decay constraints, since these depend sensitively on post-GUT physics, about which little is known. ",
        "conclusions": "We have discussed here direct detection of Milky Way neutralinos for SUGRA type models with R-parity invariance and grand unification at the GUT scale. By combining data from a variety of sources, e.g. Higgs mass bound, $b\\rightarrow s\\gamma$ branching ratio, relic density constraints and the possible new muon magnetic moment anomaly  of the BNL E821 experiment, one can greatly sharpen predictions.  For the mSUGRA model, the $m_h$ and $b\\rightarrow s\\gamma$ bounds create a lower bound on $m_{1/2}$ of $m_{1/2} \\stackrel{>}{\\sim}(300-400)$GeV (i. e. $m_{\\tilde\\chi^0_1}  \\stackrel{>}{\\sim} (120-140)$GeV). Thus puts the parameter space mostly in the $\\tilde\\tau_1-\\tilde\\chi^0_1$ co-annihilation domain, which strongly correlates $m_0$ with $m_{1/2}$. For $\\mu > 0$ and $m_{1/2}< 1$TeV, one finds $\\sigma_{\\tilde\\chi^0_1-p}\\stackrel{<}{\\sim} 10^{-10}$ pb which is within the upper reach of future planned dark matter detectors, while for $\\mu < 0$ there will be large regions unaccessible to such detectors. If the $a_\\mu$ anomaly is confirmed, then $\\mu >0$ and $m_{1/2} < 800$ GeV.  Non-universal soft breaking models allow one to raise $\\sigma_{\\tilde\\chi^0_1-p}$ by a factor as large as 10 - 100, which could account for the large  cross sections of the DAMA data. They can also open new allowed regions of the $m_0 - m_{1/2}$ plane from the $Z$ channel annihilation in the relic density calculation. The new $Z$-channel regions have larger cross sections, though still below the DAMA region, but they should be accessible when CDMS is in the SOUDAN mine and to the GENIUS-TF detector. Thus dark matter detectors should be able to investigate the nature of SUSY soft breaking, i.e. the nature of the post-GUT physics that determine the soft breaking pattern."
    },
    "0201/astro-ph0201114_arXiv.txt": {
        "abstract": "{We present Newtonian three--dimensional hydrodynamical simulations of the merger of quark stars with black holes. The initial conditions correspond to non-spinning stars in Keplerian orbits, the code includes gravitational radiation reaction in the quadrupole approximation for point masses. We find that the quark star is disrupted, forming transient accretion structures around the black hole, but 0.03 of the original stellar mass survives the initial encounter and remains in an elongated orbit as a rapidly rotating quark starlet, in all cases. No resolvable amount of mass is dynamically ejected during the encounters---the black hole eventually accretes $99.99\\%\\pm0.01\\%$ of the quark matter initially present.} ",
        "introduction": "In this paper we study the binary coalescence of a black hole and a quark star. Stellar population studies indicate that if quark stars and black holes exist at all, such binaries should exist in numbers significant from the point of view of next-generation laser interferometric gravitational wave detectors, but smaller than the number of Hulse-Taylor type binaries (Belczy\\'nski et al. 2001). Coalescing quark stars also remain strong candidates for gamma-ray burst sources (Paczy\\'nski 1991, 2001; Haensel et al. 1991). This is reason enough to study such coalescences. However, in this numerical study we address only one specific question: how much quark matter, and with what velocities, is ejected when a quark star coalesces with a black hole. The interest here is in the speculations that such ejecta may convert all neutron stars to quark stars. ",
        "conclusions": "We have found no convincing evidence of ejection of quark matter from the binaries modeled. Further simulations, not reported here, show that inclusion of pseudo-Newtonian potentials only strengthens this conclusion. This may encourage other workers to continue studying the astrophysics of quark stars."
    },
    "0201/astro-ph0201195.txt": {
        "abstract": "We present numerical magnetohydrodynamic (MHD) simulations of the effect of stellar dipole magnetic fields on line-driven wind outflows from hot, luminous stars. Unlike previous fixed-field analyses, the simulations here take full account of the dynamical competition between field and flow, and thus apply to a full range of magnetic field strength, and within both closed and open magnetic topologies. A key result is that the overall degree to which the wind is influenced by the field depends largely on a single, dimensionless, `wind magnetic confinement parameter', $\\eta_{\\ast}$ ($ = B_{eq}^2 R_\\ast^2/{\\dot M} v_\\infty$), which characterizes the ratio between magnetic field energy density and kinetic energy density of the wind. For weak confinement $\\eta_{\\ast} \\le 1$, the field is fully opened by the wind outflow, but nonetheless for confinements as small as $\\eta_{\\ast}=1/10$ can have a significant back-influence in enhancing the density and reducing the flow speed near the magnetic equator. For stronger confinement $\\eta_{\\ast} > 1$, the magnetic field remains closed over a limited range of latitude and height about the equatorial surface, but eventually is opened into a nearly radial configuration at large radii. Within closed loops, the flow is channeled toward loop tops into shock collisions that are strong enough to produce hard X-rays, with the stagnated material then pulled by gravity back onto the star in quite complex and variable inflow patterns. Within open field flow, the equatorial channeling leads to oblique shocks that are again strong enough to produce X-rays, and also lead to a thin, dense, slowly outflowing `disk' at the magnetic equator. The polar flow is characterized by a faster-than-radial expansion that is more gradual than anticipated in previous 1D flow-tube analyses, and leads to a much more modest increase in terminal speed ($ < 30\\%$), consistent with observational constraints. Overall, the results here provide a dynamical groundwork for interpreting many types of observations -- e.g., UV line profile variability; red-shifted absorption or emission features; enhanced density-squared emission; X-ray emission -- that might be associated with perturbation of hot-star winds by surface magnetic fields. ",
        "introduction": "\\label{sec:introduction}  Hot, luminous, OB-type stars have strong stellar winds, with asymptotic flow speeds up to $v_\\infty \\sim 3000$~km/s and mass loss rates up to ${\\dot M} \\sim 10^{-5} M_\\odot$~/yr. These general properties are well-explained by modern extensions %\\citet[e.g.]{PPK86} (e.g. Pauldrach, Puls, and Kudritzki 1986)  of the basic formalism developed by %\\citet[hereafter CAK]{CAK} Castor, Abbott, and Klein (1975; hereafter CAK) for wind driving by scattering of the star's continuum radiation in a large ensemble of spectral lines.  However there is also extensive evidence that such winds are not the steady, smooth outflows envisioned in these spherically symmetric, time-independent, CAK-type models, but instead have extensive structure and variability on a range of spatial and temporal scales. Relatively small-scale, stochastic structure -- e.g. as evidenced by often quite constant soft X-ray emission (Long \\& White 1980), or by UV lines with extended black troughs understood to be a signature of a nonmonotonic velocity field (Lucy 1982) -- seems most likely a natural result of the strong, intrinsic instability of the line-driving mechanism itself (Owocki 1994; Feldmeier 1995). But larger-scale structure -- e.g. as evidence by explicit UV line profile variability in even low signal-to-noise IUE spectra (Kaper et al. 1996; Howarth \\& Smith 1995) -- seems instead likely to be the consequence of wind perturbation by processes occurring in the underlying star. For example, the photospheric spectra of many hot stars show evidence of radial and/or non-radial pulsation, and in a few cases there is evidence linking this with observed variability in UV wind lines (Telting, Aerts, \\& Mathias 1997; Mathias et al. 2001).  An alternate scenario -- one explored through dynamical simulations here -- is that, in at least some hot stars, surface magnetic fields could perturb, and perhaps even channel, the wind outflow, leading to rotational modulation of wind structure that is diagnosed in UV line profiles, and perhaps even to magnetically confined wind-shocks with velocities sufficient to produce the relatively hard X-ray emission seen in some hot-stars.  The sun provides a vivid example of how a stellar wind can be substantially influenced by a surface magnetic field. Both white-light and X-ray pictures show the solar corona to be highly structured, with dense loops where the magnetic field confines the coronal gas, and low-density coronal holes where the more radial magnetic field allows a high-speed, pressure-driven, coronal outflow (Zirker 1977). %GIVE SOME REFS HERE TO, E.G., SKYLAB, SOHO, TRACE, ULYSSES. In a seminal paper, Pneuman and Kopp (1971) provided the first magnetohydrodynamical (MHD) model of this competition between magnetic confinement and coronal expansion. Using an iterative scheme to solve the relevant partial differential equations for field and flow, they showed that this competition leads naturally to the commonly observed `helmet' streamer %(named after the top-pointed, Prussian army helmet of the first world war) configuration, for which the field above closed magnetic loops is extended radially outward by the wind outflow. % Nowadays such MHD processes can be readily modelled using time-dependent MHD simulation codes, such as the Versatile Advection Code (Keppens and Goedbloed 1999), or the publicly available ZEUS codes (Stone and Norman 1992). Here we apply the latter to study MHD processes within {\\it line-driven} stellar winds that have many characteristics quite distinct from the {\\it pressure-driven} solar wind.  For the solar wind, the acceleration to supersonic speeds can take several solar radii; as such, magnetic loops that typically close within a solar radius or so can generally maintain the coronal gas in a nearly hydrostatic configuration. As we show below, in the more rapid line-acceleration of hot-stars winds, strong magnetic confinement typically channels an already supersonic outflow, often leading to strong shocks where material originating from different footpoints is forced to collide, with compressed material generally falling back to the star in quite complex and chaotic patterns. In the solar wind, the very low mass-loss-rate, and thus the low gas density and pressure, mean that only a modest magnetic field strength, of order of a Gauss, is sufficient to cause significant confinement and channeling of the coronal expansion. In hot-star winds, magnetic confinement or channeling generally requires a much stronger magnetic field, on the order of hundreds of Gauss. As such, a key issue underlying the study here regards the theoretical prospects and observational evidence for hot-star magnetic fields of this magnitude.  In the sun and other cool stars, magnetic fields are understood to be generated through a dynamo mechanism, in which coriolis forces associated with stellar rotation deflect convective motions in the hydrogen and helium recombination zones. In hot stars, hydrogen remains fully ionized even through the atmosphere, and so, lacking the strong convection zones associated with hydrogren recombination, such stars have traditionally been considered not to have strong, dynamo-generated magnetic fields. However, considering the generally quite rapid rotation of most hot stars, dynamo-generation may still be possible, e.g. within thin, weaker, near-surface convection zones associated with recombination of fully ionized helium.  Moreover, the interior, energy-generation cores of such massive stars are thought to have strong convection, and recently Cassinelli and Macgregor (2000; see also Charbonneau and MacGregor 2001) have proposed that dynamo-generated magnetic flux tubes from this interior could become buoyant, and thus drive an upward diffusion to the surface over a time-scale of a few million years. Such a model would predict surface appearance of magnetic fields in hot-stars that have evolved somewhat from the zero-age-main sequence. Alternatively, magnetic fields could form from an early, convective phase during the star's initial formation, or perhaps even arise  through compression of interstellar magnetic flux during the initial stellar collapse. Such primordial models would thus predict magnetic fields to be strongest in the youngest stars, with then some gradual decay as the star evolves.  In recent years there has been considerable effort to develop new techniques (e.g., based on the Hanle effect; Ignace, Nordsieck, \\& Cassinelli 1997; Ignace, Cassinelli, \\& Nordsieck 1999) to observationally detect stellar magnetic fields. But the most direct and well-demonstrated method is through the Zeeman splitting and associated circular polarization of stellar photospheric absorption lines (Borra \\& Landstreet 1980). This technique has been used extensively in direct measurement of the quite strong magnetic fields that occur in the chemically peculiar Ap and Bp stars (Babcock 1960; Borra et al. 1980; Bohlender 1993; Mathys 1995; Mathys et al. 1997). For more `normal' (i.e., chemically non-peculiar) hot stars, the generally strong rotational line-broadening severely hinders the direct spectropolarimetric detection of their generally much weaker fields, yielding instead mostly only upper limits, typically of order a few hundred Gauss. This, coincidentally and quite tantalizingly, is similar to the level at which magnetic fields can be expected to become dynamically significant for channeling the wind  outflow.  Recently, however, there have been first reports of positive field detections in a few normal hot stars. For the relatively slowly rotating star $\\beta$~Cephei, Henrichs et al. (2000) and Donati et al. (2001) report a 3-sigma detection of a ca. 400 G dipole field, with moreover a rotational modulation suggesting the magnetic axis is tilted to be nearly perpendicular to the stellar rotation. There are also initial reports (Donati 2001) of a ca. 1000 G dipole field in $\\theta^1$~Ori~C, this time with the magnetic axis tilted by about 45 degrees to the rotation. %\\par For Ap and Bp stars, the most generally favored explanation for their strong fields is that they may be primordial, and the relative youth of $\\theta^1$~Ori~C also seems to suggest a primordial model. In contrast, the more evolved evolutionary status of $\\beta$~Cephei seems to favor the interior-eruption scenario. \\par The focus of the present paper is to carry out magnetohydrodynamical simulations of how such magnetic fields on the surface of hot stars can influence their radiatively driven stellar wind. Our approach here represents a natural extension of the previous studies by Babel \\& Montmerle (1997a,b; hereafter BM97a,b), which effectively {\\it prescribed} a fixed magnetic field geometry to channel the wind outflow. For large magnetic loops, wind material from opposite footpoints is accelerated to a substantial fraction of the wind terminal speed (i.e. $\\sim 1000$~km/s) before the channeling toward the loop tops forces a collision with very strong shocks, thereby heating the gas to temperatures ($10^7-10^8$~K) that are high enough to emit hard (few keV) X-rays. This `magnetically confined wind shock' (MCWS) model was initially used to explain X-ray emission from the Ap-Bp star IQ Aur (BM97a), which has a quite strong magnetic field ($\\sim 4$kG) and a rather weak wind (mass loss rate $\\sim 10^{-10} M_{\\odot}/$~yr), and thus can indeed be reasonably modeled within the framework of prescribed magnetic field geometry.\\footnote{However, note that even in this case the more-rapid radial decline in magnetic vs. wind energy density means that the wind outflow eventually wins, drawing out portions of the surface field into a radial, open configuration. Such open-field regions can only be heuristically accounted for in the fixed-field modeling approach.} Later, BM97b applied this model to explain the periodic variation of X-ray emission of the O7 star $\\theta^1$~Ori~C, which has a much lower magnetic field ($\\ltwig 1000$ G) and significantly stronger wind (mass loss rate $\\sim 10^{-7} M_{\\odot}/$~yr), raising now the possibilty that the wind itself could influence the field geometry in a way that is not considered in the simple fixed-field approach. \\par The simulation models here are based on an isothermal approximation of the complex energy balance, and so can provide only a rough estimate of the level of shock heating and X-ray generation. But a key advantage over previous approaches is that these models do allow for such a fully dynamical competition between the field and flow. A central result is that the overall effectiveness of magnetic field in channeling the wind outflow can be well characterized in terms of single `wind magnetic confinement parameter' $\\eta_\\ast$, defined in eqn. (\\ref{wmcpdef}) below,  and related to the relative energy densities of field and wind (\\S 3). The specifics of our numerical MHD method are described in \\S 2, while \\S 4 details the results of a general parameter study of hot-star winds with various degrees of magetic confinement. Following a discussion (\\S 5) of the implications of these results for modeling hot-star wind structure and variability, we finally conclude (\\S 6) with a summary  and outlook for future work. ",
        "conclusions": "\\subsection{Comparison of MHD Simulations with Heuristic Scaling Estimates}  The above results lend strong support to the general idea, outlined in \\S 3.1, that the overall effect of a magnetic field in channeling and confining the wind outflow depends largely on the single magnetic confinement parameter $\\eta_{\\ast}$. Let us now consider how well these MHD simulation results correspond to the heuristic estimates for the Alfven radius $R_A$ and magnetic closure colatitude $\\theta_A$ defined in \\S 3.2.  Figure \\ref{fig6} plots contours of the Alfven radius obtained in the numerical MHD simulations with various $\\eta_{\\ast}$. Reflecting the stronger field and so higher Alfven speed, the models with larger confinement parameter have a higher Alfven radius. Note, moreover, that for all cases the Alfven radius generally decreases toward the equator. In part, this just reflects the Alfven speed associated with the dipole surface magnetic field, which has a lower strength near the magnetic equator.  But the comparison in figure \\ref{fig6} shows a systematic discrepancy between the curves showing the expected Alfven radius from this dipole model and the points showing the actual MHD results. Specifically, the dipole model underestimates the MHD Alfven radius over the pole, and overestimates it at the equator.  For the polar wind, this can be understood as a consequence of the radial stretching of the field. Figure \\ref{fig7} plots the radial variation of the polar field ratio \\begin{equation} f_{pole} (r) \\equiv { R_\\ast^2 %B (R_\\ast,\\theta=0) B_o \\over r^2 %B(r,\\theta=0) B(r,0) } \\label{fpoledef} \\end{equation} for the various magnetic confinement parameters $\\eta_{\\ast}$. For comparison, a dipole field (with $B \\sim r^{-q}$ and $q=3$) would just give a straight line of unit slope (dashed line), whereas a pure monopole, radial field (with $q=2$) would give a horizontal line at value unity, $f_{pole}=1$.  % eta        fm          sig % 10         2.43        1.13 % s10        1.98         .98 % 1          1.73         .89 % 1/s10      1.44         .79 % 1/10       1.23         .73 %  The results show that the MHD cases are intermediate between these two limits. For the weakest confinement $\\eta_{\\ast} = 1/10$, the curve bends toward the horizontal at quite small heights, reflecting how even the inner wind is strong enough to extend the polar field into a nearly radial orientation and divergence. For the strongest confinement $\\eta_{\\ast} = 10$, the field divergence initially nearly follows the dashed line for a dipole ($q=3$), but then eventually also bends over as the wind ram pressure overwhelms the magnetic confinement and again stretches the field into a nearly radial divergence. The intermediate cases show appropriately intermediate trends, but in all cases it is significant that the radial decline in field strength is generally less steep than for a pure dipole, i.e. $q<3$. The dotted line in figure \\ref{fig1} indeed shows that the MHD results for the polar Alfven radii of the various confinement cases are in much better agreement with a simple scaling that assumes a radial decline (power index $q=2.6$) that is intermediate between the dipole ($q=3$) and monopole (i.e. radial divergence, $q=2$) limits.  Overall then, at the poles the radial stretching of field by the outflowing wind has the net effect of reducing the radial decline of field, and thus increasing the Alfven radius over the value expected from the simple dipole estimate of eqn. (\\ref{radef}).  By contrast, at the equator this radial stretching has a somewhat opposite effect, tending to {\\it remove} the predominantly {\\it horizontal} components of the equatorial dipole field, and thus leading to a lower equatorial field strength and so also a lower assoicated Alfven radius, relative to the simple dipole analysis of \\S 4.2. For example, for the lowest confinement case $\\eta_{\\ast}=1/10$, the field is extended into a nearly radial configuration almost right from the stellar surface, as shown by the top left panel of figure \\ref{fig3}; the equatorial polarity switch of this radial field thus implies a vanishing equatorial Alfven speed, which thus means that contours of Alfven radius must bend sharply inward toward the surface near the equator. % For the strongest confinement case $\\eta_{\\ast}=10$, the near-surface horizontal field within closed magnetic loops about the equator remains strong enough to resist this radial stretching by wind outflow; but the faster radial fall-off in magnetic vs. flow energy means that the field above  these closed loops is eventually stretched outward into a radial configuration, thus again leading to a vanishing equatorial field and an associated inward dip in the Alfven radius.  %Overall then, at the equator the radial stretching of field by the outflowing %wind has the net effect of removing horizontal field and thus %decreasing the Alfven radius over the value expected from the simple %dipole estimate of eqn. (\\ref{radef}).  This overall dynamical lowering of the equatorial strength of magnetic field further means that the latitudinal extents of closed loops in full MHD models are generally below what is predicted by the simple dipole estimate of eqn. (\\ref{tcdef}). Thus, in previous semi-analytic models of BM97b, which effectively assume this type of dipole scaling, a somewhat larger surface field is needed to give the assumed overall extent of magnetic confinement.  \\subsection{Effect of Magnetic Field on Mass Flux and Flow Speed}  Two key general properties of spherical, non-magnetic stellar winds are the mass loss rate ${\\dot M} \\equiv 4 \\pi \\rho v_r r^2$ and terminal flow speed $v_\\infty$. To illustrate how a stellar magnetic field can alter these properties for a line-driven wind, figure \\ref{fig8} shows the outer boundary ($r=R_{max}$) values of the radial velocity, $v_r(R_{max},\\theta)$, and radial mass flux density, $\\rho(R_{max},\\theta) v_r(R_{max},\\theta)$, normalized by the values (given in Table 1) for the non-magnetic, spherically symmetric wind case, and plotted as a function of $\\mu=\\cos(\\theta)$ for each of our simulation models with various confinement parameters $\\eta_{\\ast}$.  There are several noteworthy features of these plots. Focussing first on the mass flux, note that in all models the tendency of the field to divert flow toward the magnetic equator leads to a general increase in mass flux there, with this equatorial compression becoming narrower with increasing field strength, until, for the strongest field, it forms the spike associated with an equatorial disk. This higher equatorial mass flux is associated with a higher density, since the equatorial flow speeds are always lower, quite markedly so for the dense, slowly outflowing disk of the strong field case.  Table 1 lists the overall mass loss rates, obtained by integration of these curves over the full range $-1<\\mu<1$. For the strong field case, the mass loss is reduced relative to the non-magnetic ${\\dot M}$, generally because the magnetic confinement and tilt of the inner wind outflow has effectively inhibited some of the base mass flux. Curiously, for the weakest magnetic confinement case $\\eta_{\\ast}=1/10$, there is actually a modest overall {\\it increase} in the mass loss. The reasons for this are not apparent, and will require further investigation.  \\subsubsection{Role of Rapid Areal Divergence in Enhancing Wind Flow Speed}  The right panel in figure \\ref{fig8} shows that the wind flow speed over the poles is enhanced relative to a spherical wind. This is quite reminiscent of the high-speed polar flow in the solar wind (Smith, Balogh, Forsyth, \\& McComas 2001; Horbury \\& Balogh 2001), which is generally understood to emanate from polar coronal holes of open magnetic field (Zirker 1977). Modeling of such high-speed solar wind has emphasized the important role of faster-than-radial area divergence in such open field regions (Kopp and Holzer 1976; Holzer 1977; Wang and Sheeley 1990).  Indeed, based on this solar analogy, MacGregor (1988) analyzed the effect of such rapid divergence on a line-driven stellar wind, assuming a simple 1D, radially oriented flow tube, as expected near the polar axis of an open magnetic field. He concluded that, because the line-driving acceleration scales inversely with density [$g_{lines} \\sim 1/\\rho^{\\alpha}$; see eqn. (\\ref{glines})], the lower density associated with faster divergence would lead to substantially faster terminal speeds, up to a {\\it factor three faster} than in a spherical wind, for quite reasonable values of the assumed flow divergence parameters. By comparison, the polar flow speed increases found in our full MHD models here are much  more modest, about 30\\% in even the strongest field case, $\\eta_{\\ast}=10$.  Because field and flow lines are locked together in the ideal MHD cases assumed here, the quantity $f_{pole}$, defined in eqn. (\\ref{fpoledef}) and plotted in figure \\ref{fig7}, actually also represents just this non-radial flow divergence factor for the polar flow. It is thus worthwhile to compare these dynamically computed divergence factors to the divergence assumed by MacGregor (1988), which was based on a heuristic form introduced originally by Kopp and Holzer (1976) for the solar case \\begin{equation} f(r) = { f_{max} \\exp[(r-R_1)/\\sigma]+ 1 - (f_{max}-1) \\exp[(R_\\ast-R_1)/\\sigma] \\over \\exp[(r-R_1)/\\sigma] + 1 } \\, . \\label{fkhdef} \\end{equation} Specifically, these previous analyses generally assumed that the rapid divergence would be confined to a quite narrow range of radius ($\\sigma = 0.1 R_\\ast$) centered on some radius ($R_1 = 1.25-2.5 \\, R_\\ast$) distinctly above the stellar surface radius $R_\\ast$. By comparison, our dynamical simulations indicate the divergence is generally most-rapid right at the wind base (implying $R_1 \\approx R_\\ast$), and extends over a quite large radial range (i.e., $\\sigma > 1 R_\\ast$). On the other hand, the MacGregor (1988) assumed values of the asymptoptic net divergence, $f_{max}= 1.25-2.0$, are quite comparable the divergence factors found at the outer boundary of our MHD simulation models, $f_{pole} (R_{max}) = 1.25-2.5$.  %As we shall detail in a follow-up paper (Owocki and ud-Doula 2002) %that will apply an extended, 1D-flow-tube analysis to the numerical %MHD simulations obtained here, These detailed differences in radial divergence do have some effect on the overall wind acceleration, and thus on the asymptotic flow speed. But it appears that the key reason behind the MacGregor (1988) prediction of a very strong speed enhancement was the neglect there of the finite-disk correction factor for the line-force (Friend and Abbott 1986; Pauldrach, Puls, and Kudritzki 1986). % With this factor included, and using the MHD simulations to define both the divergence and radial tilt-angle of the field and flow, we find that a simple flow-tube analysis is able to explain quite well our numerical simulation results for not only the polar speed, but also for the latitudinal scaling of both the speed and mass flux. In particular, we find that the even stronger increase in flow speed seen at mid-latudes ($ 1/4 < |\\mu| < 3/4$) in the strongest field model ($\\eta_{\\ast}=10$) does not reflect any stronger divergence factor, but rather is largely a consequence of a {\\it reduced} base mass flux associated with a nonradial tilt of the source flow near the stellar surface. As the flow becomes nearly radial somewhat above the wind base, the lower density associated with the lower mass flux implies a stronger line-acceleration and thus a faster terminal speed along these mid-latitude flow tubes. Further details of these findings will be given in future paper. %(Owocki and ud-Doula 2002)  \\subsection{Observational Implications of these MHD Simulations}  \\subsubsection{UV Line-Profile Variability}  It is worth emphasizing here that these dynamical results for the radial flow speed have potentially important implications for interpreting the observational evidence for wind structure and variability commonly seen in UV line profiles of hot stars (e.g., the so-called Discrete Absorption Components; Henrichs et al. 1994; Prinja and Howarth 1986; Howarth and Prinja 1989). In particular, an increasingly favored paradigm is that the inferred wind structure may arise from Corotating Interaction Regions (CIRs) between fast and slow speed wind streams. This requires a base perturbation mechanism to induce latitudinal variations in wind outflow properties from the underlying, rotating star (Mullan 1984; Cranmer and Owocki 1996). Based largely on the analogy with solar wind CIRs -- for which the azimuthal  variations in speed are clearly associated with magnetic structure of the solar corona (Zirker 1977; Pizzo 1978) --, there has been a longstanding speculation that surface magnetic fields on hot stars could similarly provide the base perturbations for CIRs in line-driven stellar winds (Mullan 1984; Shore \\& Brown 1990, Donati 2001).  However, until now, one argument {\\it against} this magnetic model for hot-star-wind structure was the expectation, based largely on the Macgregor (1988) analysis, that a sufficiently strong field would likely lead to anomalously high-speed streams, in excess of 5000~km/s, representing the predicted factor of two or more enhancement above the speed for a non-magnetic wind (Owocki 1994; BM97a). %Here are all 6 citations to MacGregor 1988: %\u00ca2001ApJ...559.1094C %Charbonneau,\u00caPaul; MacGregor,\u00caKeith\u00caB. %\tMagnetic Fields in Massive Stars. I. Dynamo Models % %\u00ca1998MNRAS.301..926P %Prinja,\u00caRaman\u00caK.; Massa,\u00caDerck; Howarth,\u00caIan\u00caD.; Fullerton,\u00caAlex\u00caW. %\tRepetitive structure in the stellar wind of HD 93843: a normal O-type star % %\u00ca1997A&A...323..121B %Babel,\u00caJ.; Montmerle,\u00caT. %\tX-ray emission from Ap-Bp stars: a magnetically confined wind-shock model for IQ Aur. % %\u00ca1996ApJ...462..469C %Cranmer,\u00caSteven\u00caR.; Owocki,\u00caStanley\u00caP. %\tHydrodynamical Simulations of Corotating Interaction Regions and Discrete Absorption Components in Rotating O-Star Winds % %\u00ca1995ApJ...440..308C %Cranmer,\u00caSteven\u00caR.; Owocki,\u00caStanley\u00caP. %\tThe effect of oblateness and gravity darkening on the radiation driving in winds from rapidly rotating B stars % %\u00ca1994Ap&SS.221....3O %Owocki,\u00caStanley\u00caP. %\tTheory review: Line-driven instability and other causes of structure and variability in hot-star winds By comparison, the wind flow speeds inferred quite directly from the blue edges of strong, saturated P-Cygni absorption troughs of UV lines observed from hot stars show only a modest variation of a few hundred km/s, with essentially {\\it no evidence} for such extremely fast speeds (Prinja et al. 1998).  The full MHD results here are much more in concert with this inferred speed variation, even for the strongest field model, for which the fastest streams are not much in excess of $\\sim$3000~km/s. Moreover, in conjunction with the reduced flow speeds toward the magnetic equator, there is still quite sufficient speed contrast to yield very strong CIRs, if applied in a rotating magnetic star with some substantial tilt between magnetic and rotation axes. Through extensions of the current 2D models to a full 3D configuration, we plan in the future to carry out detailed simulations of winds from rotating hot-stars with such a tilted dipole surface field, applying these specifically toward the interpretation of observed UV line profile variability.  \\subsubsection{Infall within Confined Loops and Red-Shifted Spectral Features}  In addition to the slowly migrating discrete absorption components commonly seen in the blue absorption troughs of P-Cygni profiles of UV lines, there are also occasional occurences of {\\it redward} features in either absorption (e.g.,  in $\\tau$~Sco; Howk et al. 2000) or emission (e.g., in $\\lambda$~Eri and other Be or B supergiant stars; Peters 1986; Smith, Peters, and Grady 1991; Smith 2000; Kaufer 2000). % Myron's suggested refs: %%Peters, G. J. 1986, ApJ, 301, L61 %Smith, M. A., Peters, G. J., & Grady, C. A. 1991, ApJ,367, 302 %Smith, M. A. 1989, ApJS, 71, 357 - see Sect. Va-iii (p. 375) %Rivinius, Th et al. 1998, A&A, %Smith, M. A. 2000, (Alicante), ASP Ser No. 214, p. 292 %Smith, M. A. & Ebbets, D. 1981, ApJ, 248, 214 %Kaufer, A. 2000, (Alicante), ASP Ser No. 214, p. 37 %Campos, A. J. & Smith, M. A. 1980, ApJ, 238, 250 (see Fig 13) Within the usual context of circumstellar material that is either in an orbiting disk or an outflowing wind, such redshifted spectral features have been difficult to understand, since they require material flowing {\\it away }from the observer, either in absorption against the stellar disk, or in emission from an excess of receding material radiating from a volume not occulted by the star. In general this thus seems to require material {\\it infall} back toward the star and onto the surface. Indeed, there have been several heuristic models that have postulated such infall might result from a stagnation of the wind outflow, for example due to clumping (Howk et al. 2000), or decoupling of the driving ions (Porter and Skouza 1999).  In this context, the dynamical MHD models here seem to provide another, quite natural explanation, namely that such infall is an inevitable outcome of the trapping of wind material within close magnetic loops whenever there sufficiently strong wind magnetic confinement, $\\eta_{\\ast} \\ge 1$. In principal, such interpretations of observed red-shifts in terms of infall within closed magnetic loops could offer the possibility of a new, indirect diagnostic of stellar magnetic properties. For example, the observed redshift speed could be associated with a minimium required loop height to achieve such a speed by gravitational infall. In future studies, we thus intend to generate synthetic line absorption and emission diagnostics for these MHD confinement models, and compare these with the above cited cases exhibiting redshifted spectral features.  \\subsubsection{Effect on Density-Squared Emission}  In addition to such effects on spectral line profiles from scattering, absorption, or emission lines, the extensive compression of material seen in these MHD models should also lead to an overall enhancement of those types of emission, both in lines and continuum, for which the volume emission rate scales with the square of the density. Specific examples include line emission from both collisional excitation or recombination, or free-free continuum emission in the infrared and radio. In principle, the former might even lead to a net emission above the continuum in the hydrogen Balmer lines, and thus to formal classification as a Be star, even without the usual association of an orbiting circumstellar disk. Such a mechanism may in fact be the origin for the occasional occurence of hydrogen line emission in slowly rotating B stars, most notably $\\beta$~Ceph, for which there has indeed now been a positive detection of a tilted dipole field of polar magnetic around 300 G (Donati et al. 2001; Henrichs et al. 2000). Again, further work will be needed to apply the dynamical MHD models here toward interpretation of observations of density-square emissions from hot stars that seem likely candidates for substantial wind magnetic confinement.  \\subsubsection{Implications for X-ray Emission}  Particularly noteworthly among the potential observational consequences of these MHD models are the clear implications for interpreting the detection of sometimes quite hard, and even cyclically variable, X-ray emission from some hot stars. As noted in the introduction, there have already been quite extensive efforts to model such X-ray emission within the context of a fixed magnetic field that channels wind flow into strong shock collisions (BM97a,b). In contrast, while the isothermal MHD models here do not yet include the detailed energy balance treatment necessary for quantitative modeling of such shocked-gas X-ray emission, they do provide a much more complete and dynamically consistent picture of the field and flow configuration associated with such magnetic channeling and shock compression.  Indeed, as a prelude to future quantitative models with explicit computations of the energy balance and X-ray emission, let us briefly apply here an approximate analysis of our model results that can yield rough estimates for the expected level of compressional heating and associated X-ray production. % The central idea is to assume that, within the context of the present isothermal models, any compressive heating that occurs is quickly balanced by radiative losses within a narrow, unresolved cooling layer (Castor 1987; Feldmeier et al. 1997; Cooper 1994). For  shock-type compressions with a sufficiently strong velocity jump, this radiative emission should include a substantial component in the X-ray bandpass.  Applying this perspective, we first identify within our simulation models locations of locally strong compressions, i.e. where there are substantial zone to zone decreases in flow speed along the direction of the flow itself. Taking into account that the quadratic viscosity within the Zeus code typically spreads any shocks over about 3 or 4 zones, we can use this to estimate an associated total shock jump in the specific kinetic energy $\\Delta v^2/2$. We then apply the standard shock jump conditions to obtain a correponding estimate of post-shock temperatures (BM97b), \\begin{equation} T_s \\approx 2.7 \\times 10^5 \\, K ~ { - \\Delta v^2 /2 \\over (100 \\, km/s)^2 } \\, . \\label{tsdef} \\end{equation}  Figure \\ref{fig9}c shows contours of $T_s$ computed in this way for the strong confinement case $\\eta_{\\ast}=10$. Note that quite high temperatures, in excess of $10^7$~K, occur in both closed loops near the surface, as well as for the open-field, equatorial disk outflow in the outer wind. For the closed loops, where the field forces material into particularly strong, nearly head-on, shock collisions, this is as expected from previous fixed-field models (BM97a,b).  But for the open-field, equatorial disk outflow, the high-temperature compression is quite unexpected. Since the flow impingent onto the disk has a quite oblique angle, dissipation of just the normal component of velocity would not give a very strong shock compression. But this point of view assumes a ``free-slip'' post-shock flow, i.e., that the fast radial flow speed would remain unchanged by the shock. However, our simulations show that the radial speed within the disk is much slower. Thus, under the more realistic assumption that incoming material becomes fully entrained with the disk material, i.e. follows instead a ``no-slip'' condition, then the reduction from the fast radial wind speed implies a strong dissipation of radial flow kinetic energy, and thus a quite high post-shock temperature.  To estimate the associated magnitude of expected X-ray emission, we first compute the local volume rate of compressive heating, obtained from the negative divergence of the local kinetic energy flux, \\begin{equation} q \\equiv - \\nabla \\cdot ( {\\bf v} \\rho v^2/2 ) \\approx - \\rho {\\bf v} \\cdot \\nabla v^2/2 \\, . \\label{qdef} \\end{equation} The contours of $q$ plotted in figure \\ref{fig9}d again show that strong compressions are concentrated toward the magnetic equator, with again substantial levels occuring in both the inner, closed loops, as well as in the equatorial disk outflow.  Let us next combine these results to estimate the X-ray emission above some minimum threshold energy $E$, weighting the emission by a ``Boltzmann factor'' that declines exponentially with the ratio of this energy to the shock energy $kT_s$, \\begin{equation} q_E \\equiv q \\,  e^{-E/kT_s} \\, . \\label{qedef} \\end{equation} Using the conversion that a soft X-ray energy threshold of $E=0.1$~keV corresponds roughly to a temperature of $1.1 \\times 10^6$~K, the contours in figure \\ref{fig9}e show that soft X-rays above this energy would again be produced in both the inner and outer regions of the equatorial disk.  Figures \\ref{fig9} d and e show that the {\\it volume} for flow compression and associated X-ray emission is quite limited, confined to narrow disk about the magnetic equator. Nonetheless, the strength of this emission can be quite significant. For example, volume integration of the regions defined in figure \\ref{fig9}d give a total rate of energy compression $L_c \\sim 10^{36}$~erg/s, which represents about 25\\% of the total wind kinetic energy, $L_w \\sim {\\dot M} v_\\infty^2/2 \\sim 4 \\times 10^{36}$~erg/s.  This is consistent with the fraction of mass loss in the slowly outflowing, equatorial disk, which has a value ${\\dot M_{eq}} \\sim 5  \\times 10^{-7} ~ M_{\\odot}$/yr, or about the same 25\\% of the total wind mass loss rate ${\\dot M} \\sim 2 \\times  10^{-6} \\, M_{\\odot}$/yr. The terminal speed within this disk, $v_{eq} \\sim$~1000~km/s, is about a third of that in the wind, $v_{\\infty} \\sim$~3000~km/s, implying nearly an order magnitude lower specific kinetic energy. The  `missing' energy associated with this slow disk outflow thus represents roughly the total wind flow kinetic energy dissipated by the flow into this slow disk.  Finally, integration of the Boltzmann-weighted emission in figure \\ref{fig9}e gives an estimate for the soft X-ray emission above 0.1~keV of $L_x \\sim 10^{35}$~erg/s. This is substantially higher than the canonical X-ray emission associated with intrinsic wind instabilities, $L_x \\sim 10^{-7} L_{bol} \\sim 4 \\times 10^{32}$~erg/s. This supports the general notion that hot-stars with anomalously large, observed X-ray luminosities might indeed be explained by flow compressions associated with wind-magnetic channeling.  While this analysis thus provides a rough estimate of the X-ray emission properties expected from such MHD models of wind magnetic confinement, we again emphasize that quantitative calculations of expected X-ray emission levels and spectra will require a future, explicit treatment of the wind energy balance."
    },
    "0201/astro-ph0201322_arXiv.txt": {
        "abstract": "We calculate Gamma-Ray Burst (GRB) afterglow light-curves from a relativistic jet of initial opening angle $\\theta_0$, as seen by observers at a wide range of viewing angles, $\\theta_{\\rm obs}$, from the jet axis. We describe three increasingly more realistic models and compare the resulting light-curves.  An observer at $\\theta_{\\rm obs}<\\theta_0$ should see a light curve very similar to that for an on-axis observer. An observer at $\\theta_{\\rm obs}>\\theta_0$ should see a rising light curve at early times, the flux peaking when the jet Lorentz factor $\\sim 1/\\theta_{\\rm obs}$. After this time the flux is not very different from that seen by an on-axis observer. A strong linear polarization ($\\lesssim 40\\%$) may occur near the peak in the light curve, and slowly decay with time. We show that if GRB jets have a universal energy, then orphan afterglows associated with off-axis jets should be seen up to a constant $\\theta_{\\rm obs}$, therefore the detection rate of orphan afterglows would be proportional to the true GRB rate. We also discuss the proposed connection between supernova 1998bw and GRB 980425. ",
        "introduction": "\\label{sec:intro}  Gamma-Ray Bursts (GRBs) are explosions which release roughly 10$^{51}$ erg in the form of kinetic energy of highly relativistic material (Frail et al. 2001, Panaitescu \\& Kumar 2001)\\footnote{ Most of the information we have about GRB explosions is only for the so-called long bursts, lasting more than a few seconds}. Many GRBs appear to be highly non-spherical explosions, as evidenced by a nearly-achromatic break in the light-curve (e.g. Harrison et al. 1999; Stanek et al. 1999). Highly relativistic jets are ``visible'' when our line of sight is within the jet aperture ($\\theta_{\\rm obs}<\\theta_0$), otherwise, because of relativistic beaming of photons away from our line-of-sight, the object is too dim. As the jet decelerates, the relativistic beaming becomes less severe and the emission from the jet becomes detectable to observers at larger viewing angles.  In this Letter we study the afterglow light-curves for off-axis locations ($\\theta_{\\rm obs}>0$), focusing on observers lying outside of the initial jet opening angle ($\\theta_{\\rm obs}>\\theta_0$). Granot et al. (2001) have shown that the light curve seen by an observer located within the initial jet aperture ($\\theta_{\\rm obs}<\\theta_0$) is very similar to that for an on-axis observer ($\\theta_{\\rm obs}=0$). Dalal et al. (2002) and Rossi et al. (2002) have presented simple models to calculate the flux in this case. We reanalyze these models in \\S2.1 and consider more realistic models in \\S2.2 \\& \\S2.3. Moderski, Sikora and Bulik (2000) have calculated off-axis light-curves with a more complex model, similar to that presented in \\S2.2.  In \\S3 we calculate the temporal evolution of the linear polarization for various $\\theta_{\\rm obs}$. In \\S4 we analyze the prospects of using the detection rate of orphan afterglows to estimate the collimation of GRB jets. In \\S5 we analyze the suggestion of Woosley, Eastman, \\& Schmidt (1999) that a relativistic jet emanating from the SN explosion and pointing away from us could explain the observations. ",
        "conclusions": "\\label{discussion}  We have presented the calculation of light-curves from a relativistic jet for an arbitrary location of the observer; much of the work in this letter is for an observer located outside the initial jet opening, $\\theta_{\\rm obs}>\\theta_0$. We have considered three different jet models of increasing sophistication; the simplest being a point source moving along the jet axis (\\S2.1), and the most sophisticated is 2D hydrodynamical simulation (\\S2.3). The basic qualitative features of the light-curves are similar in all three models, for $\\theta_{\\rm obs}>\\theta_0$. Moreover, the uniform jet model (model 2, \\S2.2) is in rough quantitative agreement with the hydro-model.  We find that \"orphan\" optical afterglows associated with off-axis jets can be observed up to a constant $\\theta_{\\rm obs}$, rather than a constant $\\theta_{\\rm obs}/\\theta_0$ as suggested by Dalal et al. (2002), if one assumes a constant energy in the jet, rather than a constant flux at the time of the jet break for an on-axis observer. This implies that future surveys for orphan afterglows may provide valuable data for the the distribution of jet opening angles $\\theta_0$ and the true event rate of GRBs. The orphan optical events discussed here can be identified from the initial rise during which the spectral slope is typically $\\beta > 0$, followed by a decay, on a time scale of $\\sim 1-30$ days, and may show a large degree of linear polarization ($\\lesssim 40\\%$). The detection of such orphan afterglows may provide a new line of evidence in favor of jetted outflows in GRBs. Recently Huang, Dai and Lu (2001) have considered another scenario (failed GRBs) for producing orphan afterglows; this would increase the detection rate of orphan afterglows. A good monitoring of optical transients may help distinguish failed GRBs from jets seen at $\\theta_{\\rm obs} > \\theta_0$, and improve our understanding of them."
    },
    "0201/astro-ph0201264_arXiv.txt": {
        "abstract": "{ We present a deep wide field H$\\alpha$ imaging survey of the central regions of the two nearby clusters of galaxies Coma and Abell~1367, taken with the WFC at the INT~2.5m telescope. We determine for the first time the Schechter parameters of the H$\\alpha$ luminosity function (LF) of cluster galaxies. The H$\\alpha$ LFs of Abell~1367 and Coma are compared with each other and with that of Virgo, estimated using the $B$ band LF by Sandage et al. (1985) and a $L$(H$\\alpha$)~$vs$~$M_{B}$ relation. Typical parameters of $\\phi^{*} \\approx 10^{0.00\\pm0.07}$~Mpc$^{-3}$, $L^{*} \\approx 10^{41.25\\pm0.05}$~erg~sec$^{-1}$ and $\\alpha \\approx -0.70\\pm0.10$ are found for the three clusters. The best fitting parameters of the cluster LFs differ from those found for field galaxies, showing flatter slopes and lower scaling luminosities $L^{*}$. Since, however, our H$\\alpha$ survey is significantly deeper than those of field galaxies, this result must be confirmed on similarly deep measurements of field galaxies. By computing the total SFR per unit volume of cluster galaxies, and taking into account the cluster density in the local Universe, we estimate that the contribution of clusters like Coma and Abell~1367 is approximately 0.25\\% of the SFR per unit volume of the local Universe. ",
        "introduction": "The strong morphology segregation observed in rich clusters of galaxies (Dressler, 1980) testifies the fundamental role played by the environment on the evolution of galaxies. Which physical mechanisms are responsible for such transformations is however still matter of debate. Several processes might alter the evolution of cluster galaxies. Some of them refer to the interaction of the galaxies with the intracluster medium (Gunn \\& Gott, 1972) and others account for the effects of gravitational interactions produced by the gravitational potential of the cluster (Merritt, 1983) or by galaxy-galaxy interactions (Moore et al. 1996, 1998, 1999). All these mechanisms can produce strong perturbations in the galaxy morphology with the formation of tidal tails, dynamical disturbances which appear as asymmetries in the rotation curves (Dale et al. 2001) and significant gas removal (Giovanelli \\& Haynes 1985; Valluri \\& Jog 1990).  Some of these processes are expected to produce changes in the star formation rates of galaxies in clusters. Several studies have addressed the issue of the influence of the cluster environment on the SFR of disk galaxies, however no agreement has been established so far: whereas some authors proposed similar or even enhanced star formation in cluster spirals than in the field (Donas et al. 1990, 1995; Moss \\& Whittle 1993, Gavazzi \\& Contursi 1994; Moss et al. 1998; Gavazzi et al. 1998; Moss \\& Whittle 2000), some others claim quenched SFRs in cluster spirals (Kennicutt 1983; Balogh et al. 1998; Hashimoto et al. 1998). This discrepancy could arise from non-uniformity of the adopted methods (UV vs. H$\\alpha$ vs. [O{\\sc ii}] data) or from real differences in the studied clusters (Virgo, Coma, Abell~1367, clusters from Las Campanas Redshift Survey, clusters at $z > 0.18$).\\\\  In particular, an enhanced fraction of spirals with circumnuclear H$\\alpha$ emission was found in the highest density regions of some nearby clusters (Moss et al. 1998; Moss \\& Whittle 2000), whereas no such difference was found for galaxies with diffuse emission. The compact H$\\alpha$ emission seems associated with ongoing interactions of galaxies, but numerical simulations by Bekki (1999) showed that mergers between clusters and subclusters might produce central starbursts in cluster spirals.  Existing studies of the H$\\alpha$ properties of galaxies in clusters suffer from various biases: the photoelectric data by Kennicutt et al. (1984) and Gavazzi et al. (1991, 1998) are based on samples of galaxies selected on the basis of their optical properties, independent of their H$\\alpha$ properties. On the other hand, the objective-prism surveys by Moss et al. (1988, 1998) and Moss \\& Whittle (2000) are H$\\alpha$ selected but they are too shallow to allow a determination of the H$\\alpha$ luminosity function as deep as desired.\\\\ With the aim of obtaining a reliable determination of the current SFR in nearby clusters of galaxies and to study the spatial distribution of the star formation regions, we undertook a deep imaging survey of a one degree $\\times$ one degree area of the Coma and Abell~1367 clusters.\\\\ Our work provides the first deep and complete study of galaxies in clusters based on their H$\\alpha$ emission properties.\\\\  This paper is arranged as follows: Section~2 contains a description of the observations, of the data reduction and the detection procedures. The H$\\alpha$ data are presented in Section~3. The H$\\alpha$ luminosity function and a brief discussion on the contribution of both clusters to the local star formation rate density are presented in Section~4. Conclusions are presented in Section~5. Comments on the most interesting objects as well as the H$\\alpha$ images of the detected galaxies are given in the Appendix.  \\begin{figure}[t] \\centering \\caption{ Transmitance of the filters used for the observations. } \\label{filters} \\end{figure} ",
        "conclusions": "We have carried out an H$\\alpha$ imaging survey of the central 1~deg$^{2}$ of the nearby clusters Abell~1367 and Coma. Significant H$\\alpha$ emission is found associated with 41 galaxies in Abell~1367 and 22 in Coma. These data are used to estimate, for the first time, the H$\\alpha$ luminosity function of 2 nearby clusters of galaxies. These LFs are found consistent with the H$\\alpha$ luminosity function derived for the Virgo cluster, despite their different nature. The typical Schechter parameters: $\\phi^{*} \\approx 10^{0.00\\pm0.07}$~Mpc$^{-3}$, $L^{*} \\approx 10^{41.25\\pm0.05}$~erg~sec$^{-1}$ and $\\alpha \\approx -0.70\\pm0.10$ are obtained.  \\begin{figure} \\centering \\caption{ $M_{B}$~$vs$~$\\log L(\\mbox{H}\\alpha)$ for a large sample of Virgo galaxies. The solid line indicates the adopted fit, the dashed line shows the best fit obtained from a least squares fitting. A distance modulus of 31.7~mag was adopted to convert observed fluxes and magnitudes to luminosities and absolute magnitudes. } \\label{virgo_rel} \\end{figure}  The best fitting parameters of the cluster LFs are significantly different from those found for field galaxies, in particular at the faint end where the cluster slope is shallower than the extrapolated slope of the field LF. However it must be stressed that the steep slope found in the field is based on relatively high luminosity points and no data are available below  $\\log L(\\mbox{H}\\alpha) \\approx 40$~erg~sec$^{-1}$ i.e. where the cluster LFs begin to flatten out. After re-normalizing the cluster data on the field ones, the two sets of data points are found consistent within the completeness limit of the field samples. Until a deeper field LF will be available it is impossible to establish whether the apparent underabundance of low luminosity objects in clusters is a real evolutionary effect or it is an artifact due to incompleteness.  By computing the total SFR per unit volume of the cluster galaxies, and taking into account the cluster density in the local Universe, we estimate that the contribution of type 2 and type 1 clusters is  about 0.25\\% and 10.8\\% respectively of the SFR per unit volume of the local Universe.  \\begin{figure} \\centering \\caption{ Same as Figure~6. We included the expected curve for the Virgo cluster assuming the $B$ band luminosity function from Sandage et al. (1985) and the $L(\\mbox{H}\\alpha)~vs~M_{B}$ relationship given by equation~5. } \\label{lf_clusters} \\end{figure}"
    },
    "0201/astro-ph0201502_arXiv.txt": {
        "abstract": "We present an extended analysis of the SCUBA observations of the Hubble Deep Field (HDF), expanding the areal coverage of the Hughes et al. 1998 study by a factor of $\\sim1.8$ and containing at least three further sources in addition to the five in that study. We also announce the public release of the reduced data products. The map is the deepest ever made in the sub-millimetre, obtained in excellent conditions (median $850\\mu$m optical depth of $0.16$). Two independent reductions were made, one with SURF and the other with a wholly algorithmic IDL analysis which we present in detail here. Of the three new sources, all appear to be at $z\\stackrel{>}{_\\sim}0.9$ and one is provisionally associated with an Extremely Red Object ($I-K>5$). There appears to be no significant cross-correlation signal between the $850\\mu$m fluctuations and sources detected by ISOCAM, the VLA or Chandra, nor with Very Red Objects ($I-K>4$), nor quasars and quasar candidates in the HDF (notwithstanding a small number of individual weak candidate detections). This is consistent with interpretations where the $850\\mu$m-selected galaxies are at higher redshifts than those currently probed by ISOCAM/VLA, and predominantly not Compton-thin AGN. There are only one or two compelling cases for the radio source being the sub-mm source. Nevertheless, most SCUBA-HDF point sources have a nearby radio source apparently well-separated from the sub-mm centroid. ",
        "introduction": "\\label{sec:introduction}  Sub-millimetre blank-field surveys represent a major time investment on the JCMT (e.g. Barger et al. 1999a,b, Lilly et al. 1999a,b, Eales et al. 1999, 2000, Fox et al. 2002, Scott et al. 2002), and will arguably be among the most important extragalactic surveys of the coming decade. Lensing cluster surveys (e.g. Smail, Ivison \\& Blain 1997) demonstrated early on the feasibility of sub-mm surveys with SCUBA (Holland et al. 1999) on the JCMT, and the extremely deep integration in the HDF by ourselves (Hughes et al. 1998) showed that sub-mm surveys are also feasible in blank fields. A total of $50$ hours were spent in the HDF, and the central $90''$ radius portion of this data has already been presented in Hughes et al. (1998); here we present the results from the remainder of the map, extending substantially into the Hubble Flanking Fields (HFF), and announce the public release of the reduced data products.  This paper is structured as follows. In Section \\ref{sec:method} we review the observing strategy and data acquisition; Section \\ref{sec:method2} discusses our data reduction algorithms and considers the source astrometry and flux calibration uncertainties (Section \\ref{sec:calibration}). Section \\ref{sec:results} presents the results of our reductions, discusses the possibility of spatially correlated noise (Section \\ref{sec:flipped}), and presents a discussion of the source extraction (Section \\ref{sec:deconvolution}). Details of the data products in public release are presented in Section \\ref{sec:release}. In section \\ref{sec:associations} we cross-identify our point sources with HDF/HFF, ISO, Chandra and VLA sources, and discuss the ambiguities with identifications of sub-mm survey point sources. We also attempt to obtain statistical sub-mm detections of sources not detected individually. Finally, in Section \\ref{sec:conclusions} we summarise our results. ",
        "conclusions": "\\label{sec:conclusions} Our publicly-available SCUBA map of the Hubble Deep Field resolves a substantial fraction of the extragalactic sub-mm background. At least half the sources appear to be at $z\\stackrel{>}{_\\sim}1$, based on our preliminary identifications. The lack of statistical cross-correlation signals with ISO, VLA or Chandra sources implies that the sources detected in these surveys are different populations and/or at different redshifts, in turn implying the sub-mm galaxies at these flux densities are mainly high-redshift ($z>1$) galaxies with bolometric luminosities dominated by star formation. We infer that the $\\mu$Jy radio population lies predominantly at $z<1$, in agreement with existing optical spectroscopy, and that the populations dominating the hard X-ray background contribute $<15\\%$ of the sub-mm extragalactic background light. Only two out of eight sub-mm sources are robustly identified with VLA sources; radio pre-selection (e.g. Barger et al. 2000, Chapman et al. 2001a,b) would therefore have recovered $\\geq2$ out of the $8$ SCUBA point sources in the HDF. Nevertheless, five of the eight sub-mm sources have radio sources much closer than would be expected by chance, but which are still not close enough to be physically identified with the sub-mm emission (provided we have not underestimated the astrometric uncertainties). Millimetre-wave interferometry has confirmed this is the case in the source HDF850.1, with a lensed near-infrared counterpart (Dunlop et al. 2002). This raises the interesting possibility that not all radio sources associatied with sub-mm galaxies are responsible for the far-infrared emission. Unambiguous identifications have so far almost exclusively been obtained using interferometric mm-wave follow-ups of brighter sub-mm sources, reinforcing the strategy of wide-area, shallow sub-mm surveys for the study of the resolved sub-mm point source population (e.g. Scott et al. 2002, Fox et al. 2002)."
    },
    "0201/astro-ph0201028_arXiv.txt": {
        "abstract": "{This paper describes the methodology currently being implemented in the EIS pipeline for analysing optical/infrared multi-colour data. The aim is to identify different classes of objects as well as possible undesirable features associated with the construction of colour catalogues.  The classification method used is based on the $\\chi^2$-fitting of template spectra to the observed SEDs, as measured through broad-band filters. Its main advantage is the simultaneous use of all colours, properly weighted by the photometric errors. In addition, it provides basic information on the properties of the classified objects (\\eg redshift, effective temperature). These characteristics make the \\xi2-technique ideal for handling large multi-band datasets. The results are compared to the more traditional colour-colour selection and, whenever possible, to model predictions. In order to identify objects with odd colours, either associated with rare populations or to possible problems in the catalogue, outliers are searched for in the multi-dimensional colour space using a nearest-neighbour criterion. Outliers with large \\xi2-values are individually inspected to further investigate their nature. The tools developed are used for a preliminary analysis of the multi-colour point source catalogue constructed from the optical/infrared imaging data obtained for the Chandra Deep Field South (CDF-S). These data are publicly available, representing the first installment of the ongoing EIS Deep Public Survey. \\keywords { Surveys - quasars - stars : white dwarves - stars: low mass, brown dwarves}}  \\date{Received 14 August 2001, accepted} \\offprints{Evanthia Hatziminaoglou, \\email{ehatzimi@eso.org}} ",
        "introduction": "\\label{intro}   The ESO Imaging Survey (EIS) is an ongoing project to carry out public imaging surveys in support of VLT.  Its primary goal is to provide multi-wavelength data sets from which samples comprising different types of extragalactic and galactic objects can be extracted for follow-up spectroscopic observations. So far the surveys conducted have used different instrument/telescope setups to carry out moderately deep observations of large areas, deep optical/infrared observations of high-galactic latitude fields, contiguous areas of the SMC and LMC and selected stellar fields including open clusters, and globular clusters (for more details see da Costa, 2001).  Altogether over 50~square~degrees of the southern sky have already been surveyed, albeit using different filter combinations and reaching different magnitudes (see the EIS web page at ``http://www.hq.eso.org/science/eis/'').  The ultimate success of these surveys will, to a large extent, depend on the ability of reliably identifying different classes of objects and extracting well-defined samples for spectroscopic follow-up observations. While colour selection is nothing new and several methods have been devised and applied in the past, the demands of modern, wide-area surveys involving large numbers of objects and passbands is relatively new and must be properly addressed. Therefore, to fully achieve the scientific goals of EIS a detailed understanding of the distribution of objects in colour space is required. This is not only necessary for the selection of spectroscopic targets but also as a verification of the colour catalogues being routinely produced by the survey pipeline.  An ideal way of tackling this problem is to combine intrinsic (\\eg spectral properties) and statistical information (\\ie spatial distribution, luminosity and mass functions, evolution) regarding different classes of known objects. The nature of objects, as characterised by its spectral properties, can be assessed by comparing the measurements obtained from the multi-colour photometry with those estimated using template spectra describing different types of objects. To take into account the statistical properties of a given population requires detailed simulations of the stellar population of our Galaxy as well as the extragalactic populations. These simulations must also satisfy observational constraints, such as sky position, completeness, photometric errors, and morphological classification. In principle, combining these two independent methods should lead to a further improvement of the classification of objects extracted from colour catalogues. As a first step towards this goal, this paper discusses the classification of the objects based exclusively on their spectral properties as derived from multi-colour observations. As a practical illustration, this analysis is applied to the multi-colour point source catalogues extracted from the recently completed optical/infrared data of the Chandra Deep Field South (CDF-S) by the EIS Deep Public Survey (DPS; Vandame et al., 2001; Arnouts et al., 2001a).  The $UBVRI$ optical data covers an area of 0.25~square~degrees. These data are complemented by $JK_s$ near-infrared observations over 0.1~square~degrees located at the centre of the area covered by the optical data. While the angular coverage is relatively small, this is the first complete data set of this survey which at the end will cover a total area of 3~square~degrees corresponding to 12 times the data presented here.  Therefore the results presented in this paper provide a first assessment of the likely outcome of this survey once completed. Using the CDF-S data as a benchmark is particularly interesting considering the large number of imaging and spectroscopic observations planned for this region, in addition to the already publicly available deep X-ray observations of Chandra (Rosati et al., 2001). These observations will provide an unprecedented multi-wavelength data set that should certainly help refine the classification algorithms being developed.  In Section~\\ref{data} the data as well as the method employed in the construction of the point source catalogue are briefly reviewed. Section~\\ref{methods} presents the methods used to classify objects based on their colour properties and to search for objects located in poorly populated regions of the colour space. These methods are applied to the CDF-S optical and optical/infrared data and the results are discussed and compared to other methods in Section~\\ref{Results}. In this section tables listing different types of objects are also presented. In Section~\\ref{evaluation} an assessment of the results is carried out by visually inspecting image cutouts and examining the photometric measurements of individual objects to evaluate the performance of catalogue production and target selection procedures. General conclusions and a discussion of the main results is presented in Section~\\ref{discussion}. Finally, in Section~\\ref{summary} a brief summary is presented. ",
        "conclusions": "\\label{discussion}  The methodology described in this paper has been developed to analyse in an objective and automatic way colour catalogues being routinely produced by the EIS pipeline in order to: assign objects to different classes of astronomical sources; allow for new discoveries; and understand the limitations of the data and the procedures adopted in the derivation of source catalogues and their colours from the association of data taken in different passbands. The ultimate goal is to define procedures to efficiently extract from imaging survey data well-defined samples, with minimal contamination, for spectroscopic follow-ups.  As a first step in this direction the method of \\xi2 fitting template spectra to the measured broad-band photometry, currently being used for estimating galaxy/quasar photometric redshifts, has been employed, extending it to include different types of galactic objects. The method is intended to replace standard classification schemes based on the analysis of one or more two-colour diagrams which becomes unmanageable for large sets of multi-band data. As currently implemented the classification scheme only considers the spectral properties of the objects, neglecting other important information as the apparent magnitude of the objects and the expected density of objects of different types (see below).  The results obtained from the automatic classification are not only consistent with those that would have been obtained from traditional methods based on two-colour diagrams but also consistent with model predictions, while minimising contamination by objects of other types.  A point worth noting is that a significant number of poor classifications stem from the fact that the passbands used in the present analysis not always provide independent information and the statistical analysis leads to artificially low significances of the resulting classifications. In order to deal with this problem other techniques which take into account the proper dimensionality of the colour space for a specific class of objects (\\eg PCA) should also be considered.  Finally, it is worth emphasising that the classifications are as good as the available spectral library. The library currently being used has been assembled from publicly available models and data and a number of classes are under-represented. Improvements in the classification method will depend on the continuous upgrade of the available spectral library. Currently, the library is being upgrade to include infrared spectra of white dwarves and low-mass stars kindly provided by S. Leggett. Adding spectra for different type of objects from other ongoing spectroscopic surveys such as the SDSS  will also be of great value in improving the current library.  In order to detect potential problems and not to overlook possible new discoveries the \\xi2-method has been complemented by a procedure of identifying outliers using as dissimilarity measures Euclidean distances in the multi-dimensional colour space and adopting a nearest-neighbour isolation criterion. Despite its simplicity the criterion adopted identified rare population of objects, objects with odd colours which could be traced either to real physical effects such as variability or to problems with their measured colour, demonstrating its usefulness in greatly reducing the number of cases, about 5\\% of the entire sample, that require a more detailed inspection. This number could be reduced even more by a further screening of the sample. As alluded to in the previous section, information about variability, if available, is of great use as it is a criterion based on angular separation and magnitude differences between a source and its nearest neighbour or to the mask automatically placed around very bright stars. This together with SExtractor de-blending flag and distance to masks placed around very bright stars should be used to discard objects which are likely to have the photometry affected by light contamination, the most common problem identified.  Overall, the present analysis suggests that derived catalogues are mostly free of problems. Visual inspection of several odd colour or extremely red objects has revealed that the most frequent problems are associated to the limitations of the de-blending algorithm; contamination by close neighbours; and, in some cases, residual cosmic rays located in poorly sampled regions of the image mosaic, with insufficient number of stacked images for a proper sigma-clipping.   In a follow-up paper the classification based on the spectral properties, as presented here, will be complemented with other statistics which further characterize the different populations of extragalactic/galactic objects using mock catalogues created from Monte Carlo simulations. This is particular important in the analysis of point-sources for which the stellar population makes an important contribution which varies according to the position of the sky observed. To account for this as well as reddening effects and the different filter sets used by the various surveys a population synthesis model has been combined with galactic structure models to simulate different observations."
    },
    "0201/astro-ph0201444_arXiv.txt": {
        "abstract": " ",
        "introduction": "Until 1996, there was little evidence that most galaxies were ``shy'', i.e. that they would hide their stars behind a veil of dust and turn red when forming stars, radiating the bulk of their luminosity in the infrared (IR) at a given epoch of their history. Ten years before, IRAS had unveiled a population of luminous IR galaxies exhibiting such a ``shy'' behavior, the so-called LIGs and ULIGs (with 12$\\geq log_{10}\\left(L_{\\rm IR}/L_{\\odot}\\right)\\geq$ 11 and $log_{10}\\left(L_{\\rm IR}/L_{\\odot}\\right)\\geq$ 12 respectively), which are responsible for the shape of the bolometric luminosity function of local galaxies above $\\sim$ 10$^{11}~L_{\\odot}$ (Sanders \\& Mirabel 1996). But integrated over the whole local luminosity function, LIGs and ULIGs only produce $\\sim$ 2\\,$\\%$ of the total integrated luminosity and overall only $\\sim$ 30\\,$\\%$ of the bolometric luminosity of local galaxies is radiated in the IR above $\\lambda \\sim$ 5\\,$\\mu$m. The discovery of an extragalactic background in the IR at least as large as the UV-optical-near IR one, the so-called cosmic infrared background (CIRB), with the COBE satellite (Puget et al. 1996, see references in Elbaz et al. 2002b) implied that shyness must have been more common among galaxies in the past than it is today. This was confirmed with the detection of an excess of faint mid IR (MIR) galaxies by ISOCAM onboard ISO (Elbaz et al. 1999), as well as in the far IR (FIR) with ISOPHOT onboard ISO (Dole et al. 2001) and in the sub-millimeter with SCUBA at the JCMT (see Smail et al. 2001). This excess is relative to expectations based on galaxies in the local universe. It implies that galaxies were more luminous in the IR regime and/or more numerous in the past (Chary \\& Elbaz 2001, Franceschini et al. 2001).  \\begin{figure} % \\centerline{\\psfig{file=elbazd.fig1.eps,width=12cm}} \\caption[]{{\\bf a)} IR luminosity (left axis) and SFR (right axis) corresponding to the sensitivity limits of ISOCAM (15\\,$\\mu$m, plain line), ISOPHOT (170\\,$\\mu$m, dotted line), SCUBA (850\\,$\\mu$m, dot-dashed line) and VLA/WSRT (21 cm, dashed line, using the radio-FIR correlation) as a function of redshift. K-correction from the library of template SEDs from Chary \\& Elbaz (2001). {\\bf b)} $L_{\\rm IR}$[8-1000\\,$\\mu$m] and SFR versus redshift and $\\lambda_{\\rm rest-frame}$ for the HDFN galaxies detected above a 15\\,$\\mu$m completeness limit of 0.1 mJy (dashed line). Filled dots: 5 AGNs (left axis only). } \\label{FIG:sfr} \\end{figure} ",
        "conclusions": ""
    },
    "0201/physics0201015_arXiv.txt": {
        "abstract": "The low-frequency electric microfield distribution in a Coulomb plasma is calculated for various plasma parameters, from weak to strong Coulomb coupling and from zero to strong electron screening. Two methods of numerical calculations are employed: the adjustable-parameter exponential approximation and the Monte Carlo simulation. The results are represented by analytic fitting formulas suitable for applications. ",
        "introduction": "Because of the Stark effect, stochastic electric microfields influence optical and thermodynamic properties of a plasma. First, they affect the profiles of spectral lines and effectively lower photoionization thresholds of atoms and ions immersed in a plasma \\cite{Griem,Fuhr}. A comparison of experimental and theoretical widths and shapes of the Stark-broadened spectral lines is widely used for plasma diagnostics (e.g., Refs.~\\cite{Jiang,Vitel}). Second, in some theoretical models of the plasma equation of state (e.g., Refs.~\\cite{HM,Nayfonov}), the microfield distribution is used in order to calculate occupation numbers of the bound species (although such a calculation is not free from principal difficulties, as discussed in Ref.~\\cite{P96}). It was shown recently \\cite{Nayfonov} that a more accurate description of the microfields entails a considerable improvement of the equation-of-state model.  In many cases, the microfield perturbation can be treated as quasistationary. Then the problem is reduced to determination of the probability distribution of the \\textit{low-frequency} component of perturbing electric fields (e.g., Ref.~\\cite{APEX2}), associated with a stochastic distribution of perturbing ions, whereas the electrons can be assumed to adjust instantaneously to a configuration of the ions. The low-frequency microfields are appropriate to use in the equation of state models \\cite{Nayfonov} and in calculation of spectroscopic line profiles for those radiative transitions whose frequency does not exceed the typical frequency of microfields produced by thermal fluctuations of the electron density. For example, Stehl\\'e and Jacquemot \\cite{SJ93} used the model microfield method to analyze the line shapes and line dissolution in hydrogen plasma spectra.  Holtsmark \\cite{Holtsmark} has derived the microfield distribution function assuming that the ions are not correlated and the electron screening is negligible. This assumption is justified for very hot or rarefied plasmas, for which the Coulomb coupling parameter \\begin{equation} \\Gamma = \\frac{(Ze)^2}{a k_B T} \\approx\\frac{1.25\\times10^4{\\rm~K}}{T}\\,n_{20}^{1/3}\\,Z^{5/3} \\end{equation} is close to zero. Here, $Ze$ is the ion charge, $T$ is the temperature, $k_B$ is the Boltzmann constant, $a=(\\frac{4\\pi}{3}n_i)^{-1/3}$ is the ion sphere radius, $n_i$ is the ion number density, and $n_{20}$ is the electron number density ($n_e=Zn_i$) in units of $10^{20}{\\rm~cm}^{-3}$. As we demonstrate below, the Holtsmark approximation is inaccurate already at $\\Gamma\\sim0.1$. In modern plasma experiments, $\\Gamma$ may approach unity, whereas in stellar matter it can be much larger. In these cases, correlations of plasma particles should not be neglected.  Various approximations were developed in the past in order to take the ion correlations into account. If $\\Gamma\\lesssim1$, one may use the methods of Baranger and Mozer \\cite{BM} or Hooper \\cite{Hooper,Hooperasymp} based on a cluster expansion in powers of density. The electron screening is usually described by a Debye-like (Yukawa) effective potential, introduced in the context of microfield distributions by Hoffman and Theimer \\cite{HoffmanTheimer}. In the limit of extremely strong coupling, $\\Gamma\\gg10$, and without screening, the harmonic oscillator model by Mayer \\cite{Mayer} is applicable, in which every ion is assumed to oscillate independently of the others around its equilibrium position at the ion-sphere center.  The first theory capable to provide reliable numerical results for strongly coupled plasmas with electron screening proved to be the adjustable-parameter exponential approximation (APEX), based on a special parametrization of the electric microfield $\\bm{E}$ produced on a selected test particle (neutral or charged ``radiator'' of charge $Z_r$) which undergoes the influence of charged plasma particles (``perturbers'' of species $\\sigma$ and of charge $Z_{\\sigma}$). This method has been developed for Coulomb systems \\cite{APEX} and adapted for screened Coulomb systems and ion mixtures \\cite{APEX-multi,APEX2}. It involves non-interacting quasiparticle representation of the electron-screened ions, designed to yield the correct second moment of the microfield distribution \\cite{Iglesias&al}: \\begin{equation} \\langle \\bm{E}\\cdot\\bm{E} \\rangle= {4\\pi n_i k_B T\\over Z_r} k_s^2 \\sum_\\sigma c_\\sigma Z_\\sigma \\int_0^\\infty dr\\,r e^{-k_s r} g_\\sigma(r), \\label{EE} \\end{equation} where $g_\\sigma(r)$ and  $c_\\sigma$ denote the radial distribution function (RDF) and the relative abundance of species $\\sigma$, respectively, and where $k_s$ is an effective electron screening wave-number. After introducing the effective single-particle field in the form \\begin{equation} \\epsilon^{*}_{\\sigma}=Z_{\\sigma} e {(1+ \\alpha_\\sigma r) \\over r^2} \\,e^{-\\alpha_{\\sigma}r}, \\end{equation} the adjustable parameters $\\{\\alpha_\\sigma \\}$ are chosen to satisfy the condition \\begin{equation} 4\\pi \\int \\epsilon^2 P_\\textrm{APEX}(\\epsilon) d\\epsilon = \\langle \\bm{E}\\cdot\\bm{E} \\rangle. \\end{equation} The expression on the left-hand side of this equation contains the parameters $\\{ \\alpha_{\\sigma } \\}$ to be determined, whereas the right-hand side can be evaluated using Eq.~(\\ref{EE}), if the RDF is known. The RDF thus provides a scheme for evaluating the APEX microfield distribution and is a central ingredient whose accuracy determines the one of the APEX microfield results. In our implementation of the APEX technique, we have used the hypernetted-chain RDF calculations \\cite{Iglesias&al,Rogers,Chabrier}.  On the other hand, with the advent of powerful computers it is now possible to calculate the microfield distribution from Monte Carlo (MC) or molecular-dynamics simulations of plasmas with the minimum of simplifying assumptions (e.g., Refs.~\\cite{AG86,Stamm,GS95,Gilles97,CG00}). Moreover, the latter methods allow one to study the effects of microfield nonuniformity \\cite{Demura,Murillo} and to simulate high-frequency microfield distributions in electron-ion plasmas (e.g., Ref.~\\cite{Filinov}). The MC technique is based on a numerical simulation of space configurations of a system of particles, whereas the molecular-dynamics technique traces the time evolution of the system. For the low-frequency microfield, dynamical effects are unimportant, and the two methods yield the same results, as demonstrated, e.g., in Ref.~\\cite{Murillo}. Therefore it is sufficient to use the MC method in this case.  With these powerful tools, the microfield distribution can be calculated now for any practically important combination of plasma parameters. However, plasma spectroscopy and equation-of-state models require knowledge of this distribution at many different points or even in continuous areas of the plasma parameter space. In this case, either extensive numerical tables or approximate analytic expressions are necessary.  We present results of calculations of the low-frequency microfield distribution function at a neutral and charged plasma point for various values of $\\Gamma$ ranging from 0 to 100 and for various values of an effective electron screening length. We consider plasmas composed of a single species of ions; in particular, in the case of a charged test particle, its charge is assumed to be equal to that of perturbers. The calculations are performed mainly by the MC method; for comparison we have done also APEX calculations. We also present analytic formulas which reproduce the calculated electric microfield probability distributions with an accuracy comparable to small differences between the MC and APEX results.  In the next section, we describe basic assumptions used in our calculations and write down some asymptotic results. In Sec.~\\ref{sect-res}, we present results of numerical calculations and analytic approximations for microfield distributions produced at a neutral or charged point by ions interacting via unscreened or screened Coulomb potentials. The results are summarized in Sec.~\\ref{sect-end}. ",
        "conclusions": "\\label{sect-end} We have calculated microfield distributions at neutral and charged test particles in a one-component plasma of ions, interacting via Coulomb potential, in various regimes from weak to strong coupling. The MC and APEX methods of calculation yield similar distributions, in agreement with previously known results \\cite{APEX}. Self-consistent elementary-function approximations for the field probability density $P(\\beta)$ and its cumulative distribution $Q(\\beta)$ are constructed in the two cases of a neutral and charged point, for a Coulomb coupling parameter $\\Gamma$ varying from 0 to $10^2$.  Furthermore, MC calculations of the microfield distribution have been performed for the screened Coulomb interaction, using the model of ions interacting via the Debye-like (Yukawa) effective potential, with an effective screening length as a second independent parameter. The dimensionless screening parameter $s$ [Eq.~(\\ref{param})] varies from 0 to 3. The whole set of numerical results for $P(\\beta)$ at various values of the coupling and screening parameters is approximated by analytic expressions.  The obtained results can be used in theoretical models of optical spectra and equations of state of Coulomb plasmas."
    },
    "0201/astro-ph0201358_arXiv.txt": {
        "abstract": "We discuss the dust distribution within photoionized regions. Assuming a geometry with a central dust cavity, which is strongly suggested by the literature, we can estimate the cavity's radius from the ratio of the infrared and radio fluxes by using a simple transfer model of Lyman continuum photons. We apply the method to a sample of the Galactic H {\\sc ii} regions. The estimated typical radius of the dust cavity of the Galactic compact H {\\sc ii} regions is about 30\\% of the Str{\\\"o}mgren radius. Taking account of uncertainties both of the observational data and the model, we can reject a dust distribution lacking a central cavity. Therefore, the dust cavity model is supported independently of the previous works. We discuss the formation mechanism of such a dust cavity and its detectability by present and future infrared facilities. ",
        "introduction": "Dust grains exist everywhere. The inside of H {\\sc ii} regions is not an exception (e.g., \\citealt{ish68, har71, mel79}). The radiation from celestial objects is always absorbed and scattered by the grains. Without correction for the dust extinction, we inevitably underestimate the intrinsic intensity of the radiation, and our understanding of the physics of these objects may be misled. In order to evaluate the amount of the dust extinction in detail, it is important to know the distribution of dust as well as its optical properties.  Theoretically, some attempts to reveal the dust distribution in H {\\sc ii} regions have been made to date \\citep{mat67,mat69,gai79a,gai79b}. According to \\cite{gai79b}, the radiation pressure on the dust grains by the central source causes a central hole of dust and gas. The dust distribution has also been observationally investigated. {}From the observed flux ratio of H $\\beta$ to continuum, \\cite{ode65} have shown that the gas-dust ratio decreases with radius in the Orion nebula. {}From radial photometric profiles at $V$-band of some H {\\sc ii} regions, \\cite{nak83} also suggest that a dust depletion zone is in the central region (See also \\citealt{yos86,yos87}). Moreover, the central dust cavity is supported by fitting to the observed infrared (IR) spectral energy distribution (SED) of H {\\sc ii} regions \\citep{chi86,chi87,chu90,fai98,gho00}. Recently, such dust geometry is preferred not only by IR SED fitting but also by the radial photometric profiles at submillimeter range \\citep{hat00}. In addition, we have some evidence of central gaseous cavities from radio observations \\citep{ter65,woo89}. Since dust and gas are well coupled each other \\citep{gai79b}, the presence of a gaseous cavity indicates a dust cavity.  Although the dust cavity is likely to exist, it has not been observed directly in the far-infrared (FIR) band, which is dominated by the thermal emission from dust. This is because the FIR observations are not advanced relative to other wavelengths, and moreover, the spatial resolving power of FIR imaging is generally weak. Recently, more detailed observational data have been obtained by {\\it ISO} (Infrared Space Observatory) and SCUBA (Submillimeter Common-User Bolometer Array), and many successive facilities are planned in the near future (SIRTF, SOFIA, ASTRO-F, ALMA, etc.). By using these new powerful facilities, we may observe the dust distribution directly. We need to estimate the detectability of the cavity by these facilities.  On the other hand, we have dveloped a method for determining the fraction of Lyman continuum (LC) photons contributing to hydrogen ionization (\\citealt{ihk01, ino01}, hereafter Paper I and Paper II, respectively). In paper I and II, we show that only half of LC photons from the central source in an H {\\sc ii} region ionizes neutral hydrogens, and the rest are absorbed by dust grains within the ionized region. If a dust cavity exists in the ionized region, however, the efficiency of the dust absorption for the LC photons becomes small. In this paper, we examine whether the estimated photon fraction is reproduced  by a model nebula with a central dust cavity.  In the next section, we describe our method for examining the dust distribution, especially the central dust cavity's radius, by using the fraction of LC photons contributing to hydrogen ionization. In section 3, we apply it to the sample H {\\sc ii} regions in the Galaxy. Then, we discuss the formation mechanism of the dust cavity and its detectability by some IR facilities in section 4. Finally, we summarize our conclusions in the last section. ",
        "conclusions": "We examine a possible distribution of dust grains in H {\\sc ii} regions. Assuming the geometry with the central dust cavity, which is suggested theoretically and observationally in the literature, we can determine the cavity's radius by using a simple transfer model of Lyman continuum photons in a dusty spherical medium. We adopt the extinction law recently developed by \\cite{wei01}, which is based on two components of grains, carbonaceous and silicate grains, and is a function of the ratio of visual extinction to color excess, $R_V$. We assume $R_V=5.5$, which is suitable for high density regions. In the transfer model, only one free parameter is included; the radius of the dust cavity.  The fraction of Lyman continuum photons contributing to hydrogen ionization (i.e. ionizing photons) is adopted as a new constraint to determine the cavity's radius. The ionizing photon fractions of 13 'compact' or 'ultra-compact' H {\\sc ii} regions in the Galaxy are determined from the ratio of infrared to radio fluxes via the method developed by us recently (Paper I and II). Its mean is 0.55. It is consistent with our previous results; a half of LC photons is absorbed by dust before they ionize the neutral hydrogen.  We examine the effect of various quantities on the relation between the ionizing photon fraction and the dust cavity's radius. Setting a fixed radius of the cavity, we find that the ionizing photon fraction decreases as the ionization parameter increases. Also we find that the fraction decreases as the value of $R_V$ decreases or as electron temperature increases.  We determine the dust cavity's radius of individual sample region so as to reproduce the fraction of ionizing photons estimated from observational data by the transfer model. The mean determined radius of the dust cavity is 0.39 times Str{\\\"o}mgren radius for all of the sample. We divide our sample into two subsamples: one is the 'ultra-compact' sample, and the other is the 'compact' sample. Our classification is based on the Str{\\\"o}mgren (or ionized) radius of the sample regions, and is consistent with that of \\cite{hab79} based on the gas densities. A typical radius of the dust cavity for the 'compact' sample is $0.30\\pm0.12$ in units of the Str{\\\"o}mgren radius, where the error includes both uncertainties of the individual estimate and the variation among sample regions. This corresponds to about 40\\% of the ionized radius and about 0.3 pc. Also, a typical cavity's radius for the 'ultra-compact' sample is $0.48\\pm0.26$ in units of the the Str{\\\"o}mgren radius, which corresponds to about 60\\% of the ionized radius and about 0.2 pc.  Since the uncertainties of these values are somewhat large and the 'ultra-compact' sample consists of only three regions, it is uncertain whether the difference of the cavity's radius between two subsamples is real and indicates an evolutionary sequence of H {\\sc ii} regions. In any case, we can reject the dust distribution without the central cavity with the significance level about 0.001\\% for both samples. Thus, we conclude that the model nebula filled with dust uniformly is invalid in order to explain the ionizing photon fraction. The dust cavity model is supported independently of the previous works.  We discuss the formation mechanism of the dust cavity. Although the radiation pressure and stellar wind by the central source can produce the central cavity, it is uncertain which is dominant mechanism. On the other hand, the dust sublimation process is not a major mechanism to form the cavity. This is because the expected cavity's radius by the sublimation process is too small to explain that obtained by us.  We discuss the detectability of the dust cavity by present and future infrared--submillimeter facilities. SIRTF is the most powerful facility to detect the cavity. Indeed, we expect that SIRTF can detect the cavities in almost all H {\\sc ii} regions of our Galaxy. SCUBA can also detect the cavities if we select nearby (typically $\\la$ 5 kpc) spherical 'compact' H {\\sc ii} regions, whose gas densities are about $10^3$ cm$^{-3}$, as the targets. Moreover, Japanese IR satellite, ASTRO-F can detect the cavities in H {\\sc ii} regions located within only 1 kpc from us. Since ASTRO-F surveys the whole of the sky, we may detect a lot of closest H {\\sc ii} regions with the central cavity."
    },
    "0201/astro-ph0201385_arXiv.txt": {
        "abstract": "{% Interpolation techniques play a central role in Astronomy, where one often needs to smooth irregularly sampled data into a smooth map. In a previous article (\\citealp{2001A&A...373..359L}, hereafter Paper~I), we have considered a widely used smoothing technique and we have evaluated the expectation value of the smoothed map under a number of natural hypotheses.  Here we proceed further on this analysis and consider the variance of the smoothed map, represented by a two-point correlation function.  We show that two main sources of noise contribute to the total error budget and we show several interesting properties of these two noise terms.  The expressions obtained are also specialized to the limiting cases of low and high densities of measurements.  A number of examples are used to show in practice some of the results obtained. ",
        "introduction": "\\label{sec:introduction}  Raw astronomical data are very often discrete, in the sense that measurements are performed along a finite number of directions on the sky.  In many cases, the discrete data are believed to be single measurements of a smooth underlying field.  In such cases, it is desirable to reconstruct the original field using interpolation techniques.  A typical example of the general situation just described is given by weak lensing mass reconstructions in clusters of galaxies. In this case, thousands of noisy estimates of the tidal field of the cluster (shear) can be obtained from the observed shapes of background galaxies whose images are distorted by the gravitational field of the cluster.  All these measurements can then be combined to produce a smooth map of the cluster shear, which in turn is subsequently converted into a projected density map of the cluster mass distribution.  One of the most widely used interpolation techniques in Astronomy is based on a weighted average.  More precisely, a positive weight function, describing the relative weight of a datum at the position $\\vec \\theta + \\vec\\phi$ on the point $\\vec\\theta$, is introduced. The weight function is often chosen to be of the form $w \\bigl( | \\vec\\phi | \\bigr)$, i.e.\\ depends only on the separation $| \\vec\\phi |$ of the two points considered.  Normally, $w$ is also a decreasing function of $| \\vec\\phi |$ in order to ensure that the largest contributions to the interpolated value at $\\vec\\theta$ comes from nearby measurements.  Then, the data are averaged using a weighted mean with the weights given by the function $w$.  More precisely, calling $\\hat f_n$ the $n$-th datum obtained at the position $\\vec\\theta_n$, the smooth map is defined as \\begin{equation} \\label{eq:1} \\tilde f(\\vec\\theta) \\equiv \\dfrac{\\sum_{n=1}^N \\hat f_n w(\\vec\\theta - \\vec\\theta_n)}{\\sum_{n=1}^N w(\\vec\\theta - \\vec\\theta_n)} \\; , \\end{equation} where $N$ is the total number of objects.  In a previous paper (\\citealp{2001A&A...373..359L}, hereafter Paper~I) we have evaluated the expectation value for this expression under the following hypothesis: \\begin{itemize} \\item The measured values $\\{ \\hat f_n \\}$ are independent random variables with expectation value \\begin{equation} \\label{eq:2} \\bigl\\langle \\hat f_n \\bigr\\rangle = f(\\vec\\theta_n) \\; . \\end{equation} In other words, the $\\{ \\hat f_n \\}$ are \\textit{unbiased\\/} measurements of a field $f(\\vec\\theta)$. \\item The positions $\\{ \\vec\\theta_n \\}$ are independent random variables with uniform distribution and density $\\rho$.  In Paper~I we initially considered a fixed number $N$ of positions inside a field $\\Omega$ of finite area $A$; then, we took the \\textit{continuous limit\\/} letting $N$ go to infinity with $\\rho = N/A$ constant.  Equivalently, we considered $N$ to be a Poisson random variable with average $\\rho A$: \\begin{equation} \\label{eq:3} p_N(N) = \\e^{-\\rho A} \\frac{(\\rho A)^N}{N!} \\; , \\end{equation} and each location $\\vec\\theta_n$ to be uniformly distributed inside $A$: \\begin{equation} \\label{eq:4} p_\\theta(\\vec\\theta_n) = \\frac{1}{A} \\; . \\end{equation} \\end{itemize} In Paper~I we have shown that \\begin{equation} \\label{eq:5} \\bigl\\langle \\tilde f(\\vec\\theta) \\bigr\\rangle = \\int f(\\vec\\theta') w_\\mathrm{eff}(\\vec\\theta - \\vec\\theta') \\, \\diff^2 \\theta' \\; . \\end{equation} Thus, $\\bigl\\langle \\tilde f \\bigr\\rangle$ is the convolution of the unknown field $f$ with an \\textit{effective weight\\/} $w_\\mathrm{eff}$ which, in general, differs from the weight function $w$.  We also have shown that $w_\\mathrm{eff}$ has a ``similar'' shape as $w$ and converges to $w$ when the object density $\\rho$ is large; however for finite $\\rho$, $w_\\mathrm{eff}$ is broader than $w$.  Here we proceed further with the statistical analysis by obtaining an expression for the two-point correlation function (covariance) of this estimator.  More precisely, given two points $\\vec\\theta_A$ and $\\vec\\theta_B$, we consider the two-point correlation function of $\\tilde f$, defined as \\begin{equation} \\label{eq:6} \\Cov(\\tilde f; \\vec\\theta_A, \\vec\\theta_B) \\equiv \\bigl\\langle \\tilde f(\\vec\\theta_A) \\tilde f(\\vec\\theta_B) \\bigr\\rangle - \\bigl\\langle \\tilde f(\\vec\\theta_A) \\bigr\\rangle \\bigl\\langle \\tilde f(\\vec\\theta_B) \\bigr\\rangle \\end{equation} In our calculations, similarly to Paper~I, we assume that $\\hat f_n$ are \\textit{unbiased and mutually independent\\/} estimates of $f(\\vec\\theta_n)$ [cf.\\ Eq.~\\eqref{eq:2}].  We also assume that the $\\{ \\hat f_n \\}$ have fixed variance $\\sigma^2$, so that \\begin{equation} \\label{eq:7} \\bigl\\langle \\bigl[ \\hat f_n - f(\\vec\\theta_n) \\bigr] \\bigl[ \\hat f_m - f(\\vec\\theta_m) \\bigr] \\bigr\\rangle = \\sigma^2 \\delta_{nm} \\; . \\end{equation}  The paper is organized as follows.  In Sect.~\\ref{sec:summary} we summarize the results obtained in this paper.  In Sect.~\\ref{sec:eval-covar} we derive the general expression for the covariance of the interpolating techniques and we show that two main noise terms contribute to the total error.  These results are then generalized in Sect.~\\ref{sec:vanishing-weights} to include the case of weight functions that are not strictly positive.  A useful expansion at high densities $\\rho$ of the covariance is obtained in Sect.~\\ref{sec:moments-expansion}.  Section~\\ref{sec:properties} is devoted to the study of several interesting properties of the expressions obtained in the paper.  Finally, in Sect.~\\ref{sec:examples} we consider three simple weight functions and derive (analytically or numerically) the covariance for these cases. Four appendixes on more technical topics complete the paper. ",
        "conclusions": "\\label{sec:conclusions}  In this article we have studied in detail the covariance of a widely used smoothing technique.  The main results obtained are summarized in the following items. \\begin{enumerate} \\item The covariance is composed of two main terms, $T_\\sigma$ and $T_\\mathrm{P}$, representing measurement errors and Poisson noise, respectively; the latter one depends on the field $f$ on which the smoothing is performed. \\item Expressions to compute $T_\\sigma$ and $T_\\mathrm{P}$ have been provided.  In particular, it has been shown that both terms can be obtained in term of a kernel $C(w_A, w_B)$, which in turn can be evaluated from the weight function $w(\\vec\\theta)$. \\item We have obtained an expansion of the kernel $C(w_A, w_B)$ valid at high densities $\\rho$. \\item We have shown that $T_\\sigma$ has an upper limit, given by $\\sigma^2$, and a lower limit, provided by Eq.~\\eqref{eq:90}. \\item We have evaluated the form of the noise contributions in the limiting cases of high and low densities. \\item We have considered three typical cases of weight functions and we have evaluated $C(w_A, w_B)$ for them. \\end{enumerate}  Finally, we note that although the smoothing technique considered in this paper is by far the most widely used in Astronomy, alternative methods are available.  A statistical characterization of these methods, using a completely different approach, will be presented in a future paper (Lombardi \\& Schneider, in preparation)."
    },
    "0201/astro-ph0201450_arXiv.txt": {
        "abstract": "Observations of galaxy-galaxy lensing from Sloan Digital Sky Survey (SDSS) are combined with the Tully-Fisher and fundamental plane relations to derive constraints on galactic halo profiles. We show that both for early and late type galaxies around $L_*$ the rotation velocity decreases significantly from its peak value at the optical radius to the virial radius $r_{200}$, $v_{\\rm opt}/v_{200} \\sim 1.8$ with about 20\\% uncertainty. Such a decrease is expected in models in which the halo profile is very concentrated, so that it declines steeper than isothermal at large radii. This large decrease can be explained as a result of both a concentrated dark matter profile and a significant stellar contribution to the rotation velocity at the optical radii. We model the stellar component with a thin rotationally supported disk or a Hernquist profile and use adiabatic dark matter response model to place limits on the halo concentration as a function of the stellar mass to light ratio. For reasonable values of the latter we find concentrations $c_{200} $ consistent with CDM predictions, suggesting  there is no evidence for low concentrations for the majority of halos in the universe. We also discuss the origin of Faber-Jackson relation $L \\propto \\sigma^4$ in light of $L\\propto v_{200}^{2.5}$ relation found for early type galaxies above $L_*$ from galaxy-galaxy lensing. This leads to a decrease in $v_{\\rm opt}/v_{200}$ with luminosity above $L_*$, so that at $7L_*$ the ratio is 1.4. This is expected from the fundamental plane relation as a result of a reduction in the baryonic contribution to the total mass at the optical radius and a decrease in optical to virial rotation velocity in dark matter profile. These results imply that relations such as Tully-Fisher and Faber-Jackson are not simply those between the mass of dark matter halo and galaxy luminosity, but are also significantly influenced by the baryonic effects on the rotation velocity at optical radii. ",
        "introduction": "The study of dark matter profiles around galaxies has been an active area since the original discovery of dark matter from the flat rotation curves in spirals (see \\citeNP{2001ARA&A..39..137S} for a review). These show that there must be dark matter in the outer parts of the galactic halos, but its extent is uncertain because of the limited range probed by observations. More recent observational studies of rotational velocity in dwarf and low surface brightness galaxies have suggested that the amount of dark matter in the central regions is smaller than predicted for the average galactic population by CDM models (e.g. \\citeNP{1994Natur.370..629M}, \\citeNP{1994ApJ...427L...1F}, \\citeNP{2000ApJ...543..704D}, \\citeNP{2001ApJ...552L..23D}).  The CDM models predict very concentrated dark matter haloes. This is usually parametrized using a universal dark matter profile such as NFW \\cite{1997ApJ...490..493N} \\footnote{Although other profiles have been proposed that differ significantly from NFW in the inner parts of the halo, they agree well with NFW in the outer parts \\shortcite{2001ApJ...554..903K}. Since for the present work we are concerned predominantly with the outer parts we only use NFW profile.}. For a given virial mass of the halo, in this paper defined as the mass enclosed within a sphere of radius $r_{200}$ within which the density is 200 times critical density, the profile has a characteristic shape which depends only on a single parameter. For NFW profile, in which the slope continuously changes from the inner value of -1 to the outer value of -3, the free parameter is often defined as the concentration $c_{200}$, which is the ratio between the radius $r_{200}$ and the scale radius $r_s$ where the slope is close to $-2$. Higher concentration parameters imply higher densities at the scale radius $r_s$. On galactic scales CDM models predict $c_{200}\\sim 7-15$ depending on the halo mass, matter density, shape and amplitude of the power spectrum.  However, observational case for low concentrations is not clear and different conclusions have been reached by other studies (e.g. \\citeNP{2001MNRAS.325.1017V}). Even if the observational results of low dark matter density in the inner regions are confirmed, they do not necessarily indicate a problem for the CDM models. The variation in the profile shapes between halos is large and it is not clear whether the dwarf or low surface brightness galaxies, which show strongest evidence for low density cores, can be associated with the average halo population. For example, if these galaxies are associated with a population that had a major merger in the recent past, this may lead to a significantly flatter and less concentrated halo structure than the average population (e.g. \\shortciteNP{2001astro.ph..8151W}). Another possibility is that astrophysical processes such as bar rotation redistribute the dark matter in the inner parts of the halo \\cite{2001astro.ph.10632W}. However, if the case of low concentrations is extended to the population of galaxies as a whole and to the outer parts of the halo, which cannot be affected by astrophysical processes, then the CDM crisis would become much more severe.  It would thus be useful to have some information on the halo structure for the mean population of galaxies. While many observationally determined rotation velocity profiles exist, the main uncertainty has always been the relative contribution between the baryonic component in the gas, bulge/ellipsoid and disk and the dark matter in the halo. The baryonic and dark matter components are difficult to separate because the conversion from the star light (and, to a lesser extent, gas density) to baryonic mass is still uncertain from the stellar population synthesis models and other studies. The situation is further complicated by the fact that baryonic component is not dynamically negligible and during its condensation it induces a response of the dark matter halo changing its original distribution in the inner parts. While these problems are ameliorated for the low surface brightness galaxies, they are almost never negligible.  Since the modelling of relative disk or spheroid and halo contributions is rather difficult in the inner parts it is useful to concentrate on the outer parts of the halo, where the baryon influence is less important. However, optical rotation curves typically extend only out to $10-20$kpc and even the ones based on HI  measurements do not extend beyond 30-50kpc. Similarly, velocity dispersion studies in early type galaxies also do not extend past a few effective radii and even X-ray studies of large ellipticals do not extend beyond 10 effective radii \\cite{1999ApJ...518...50L}. Over this limited range the data on disk $L_*$ galaxies indicate that the rotation curves rise in the inner parts of the disk and then stay approximately constant or decline somewhat out to the outer limit of observations. This flatness indicates that approximating the density profile as isothermal, $\\rho(r) \\propto r^{-2}$, is a good approximation to the matter profile over this range. On the other hand, theoretical CDM models predict that the velocity profile of the dark matter increases out to twice the scale radius $r_s$ and then slowly declines beyond that as the slope gradually decreases towards -3. The fact that we do not observe a decrease in rotation velocity at radii below $2.15r_s \\sim 30-50$kpc is presumably due to the baryonic effects, which increase the mass in the center. The decrease at large radii is however a robust prediction of these models.  While, until recently, no accurate determination of the mass profile at large radii has been available, recent galaxy-galaxy lensing observations by the SDSS team \\shortcite{2001astro.ph..8013M} have improved the situation significantly by obtaining the morphology and luminosity dependence of the signal. Theoretical analysis of lensing data must take into account not just galactic halos, but also those from groups and clusters, which dominate on large scales above 200-300$h^{-1}$kpc \\cite{gs02}. These results show that a late type $L_*$ galaxy, with $L_I=2.7\\times 10^{10}h^{-2}L_{\\sun}$ (where we have applied a 30\\% internal extinction correction in $I$ band), has a mass $M_{200}\\sim (3.4\\pm 2.1)\\times 10^{11}h^{-1}M_{\\sun}$, corresponding to the circular velocity at the virial radius $v_{200} \\sim  115$km/s. While the error is still quite large it is gaussian distributed in mass, so an increase in $v_{200}$ by 35\\% to 150km/s is excluded at 95\\% confidence level. This number can be compared to the maximum rotation velocity for such a galaxy, which is 208km/s with a small scatter \\shortcite{1997ApJ...477L...1G}, the well known Tully-Fisher relation. Comparing the two values shows that a decrease in rotation velocity from the optical to the virial radius indeed occurs for the average population of spiral galaxies. The decrease is large, almost a factor of 2 and even if we push the virial velocity up by 2-$\\sigma$ to $v_{200}= 150$km/s the decrease is still around 40\\%. As shown in this paper, a similarly large decrease in rotation velocity is obtained also for early type galaxies which are not rotationally supported.  These results are interesting, since they are in the direction predicted by the CDM models and demonstrate that the halo profiles indeed become steeper than -2 in the outer parts of the halo. They set tight limits both on the structure formation models, by limiting the acceptable range of concentration parameters, and on the disk/spheroid formation models, by constraining the stellar mass to light ratio. Only for specific values of stellar mass to light ratio and halo concentration can one satisfy these constraints. The purpose of this paper is to investigate the constraints in detail using halo and disk/spheroid formation models.  Previous work on this subject has explored the constraints from the rotation velocity at optical radius given by the zero point of Tully-Fisher relation (\\citeNP{2000MNRAS.318..163M}, \\citeNP{2000ApJ...538..477N}, \\citeNP{2001ApJ...554..114E}). In the absence of virial mass information hese models must rely on additional assumptions to derive the constraints on cosmological models. The advantage of the additional information from lensing is that it provides another dynamical constraint at large radii, which can remove some of the modelling uncertainties present in previous modelling. In addition, while previous work only explored the constraints from late type galaxies, in this paper we also investigate the constraints from the early type galaxies. We find these are more robust both because the virial masses are more accurately determined and because the velocity dispersions at optical radii are obtained from the same SDSS sample. ",
        "conclusions": "In this paper we compare average properties of rotation velocities between optical and virial radii based on optical and galaxy-galay lensing measurements, respectively. Our model is statistical in nature, since it is not based on analysis of individual galaxies, but on their average properties as a function of luminosity. On the other hand, the galaxies in our sample were not chosen on the basis of any selection criteria, so our results should apply to the galaxy population as a whole. Moreover, the large dynamical range between optical region (of order a few kpc) and virial region (of order a few hundred kpc) allows one to measure slow changes in rotation velocity which are not possible to detect from each of the observations individually.  Our main conclusion is that the rotation velocity in galaxies decreases significantly from optical radii to virial radii. Such a decrease is theoretically expected since the dark matter profiles in CDM models are steeper than isothermal at large radii and contribution from stars further increases the rotation velocity at the optical radius. It has however not been clearly demonstrated previously because of a narrow range of scales observed. This demostrates the power of g-g lensing measuring the mass at large radii combined with the more traditional methods that measure mass at smaller radii (see also \\citeNP{2001ApJ...555..572W} for a similar conclusion for early type galaxies at higher redshifts). The concentration parameters obtained using reasonable stellar mass to light ratios are around $c\\sim 10$ and are in a good agreement with predictions of $\\Lambda$CDM model with $\\Omega_m=0.3$ and $\\sigma_8=0.8$, which gives $c_{200} \\sim 10$ at $M \\sim 10^{12}h^{-1} M_{\\sun}$ \\shortcite{2001ApJ...554..114E}. We find no evidence for low halo concentrations in the main galaxy population both for early and late type galaxies.  The theoretically predicted stellar mass to light ratios have considerable theoretical uncertainties. Assuming CDM halo profiles and concentrations we find $\\Upsilon_I \\sim 1.5h$ for late type galaxies and $\\Upsilon_i \\sim 3h$ for early type galaxies. These values are in a good agreement with the stellar population synthesis models (e.g. \\citeNP{1998MNRAS.294..705K}, \\shortciteNP{2002astro.ph..1207D}). While higher stellar mass to light ratios are often quoted in the literature (e.g. \\shortciteNP{2001AJ....121.1936G}) these are based on minimal halo models and are thus often an upper limit. It is interesting to note that in our models the dark matter contribution to rotation velocity is comparable to that of stars even below effective radius. This is because of the adiabatic response, which compresses the dark matter in the center. While this model may be oversimplified it nevertheless suggests that it may be difficult or impossible to separate the two components on the basis of dynamical studies. Using the minimal halo assumption the stellar mass to light ratio in early type galaxies may be overestimated up to a factor of two.  While the optical to virial velocity ratio is above unity both for early and late type galaxies, it is decreasing with halo mass from 1.8 at $3\\times 10^{11}h^{-1}M_{\\sun}$ to 1.4 at $10^{13}h^{-1}M_{\\sun}$ (this statement is valid only for galaxies above $L_*$ and it is possible that the trend is reversed towards the low luminosity galaxies). This trend continues further into the cluster halo masses, where the ratio falls below unity. Such a trend is expected in models where halo profiles are less concentrated for higher halo masses, implying that the turnaround from an increase to a decrease in rotation velocity occurs at a larger radius relative to the virial radius. In addition, in more massive halos stars play a less important role both as a direct contribution to the rotation velocity and through their effect on the dark matter. The rotation velocity-luminosity scalings at optical radii, such as Tully-Fisher and Faber-Jackson relations, are not directly related to the properties of dark matter, but also require a proper modelling of baryons and dark matter response to baryonic contraction (see also a related discussion in \\shortciteNP{2000ApJ...528..145G}, \\citeNP{2001astro.ph..8160K} and \\citeNP{2001astro.ph.12566V}). While there are still uncertainties in the modelling of these processes, the simple models presented here reproduce well the constraints from the data both for early and late type galaxies.  How do our results compare to previous work? The zero point of TF relation problem (\\shortciteNP{2001ApJ...554..114E}, \\citeNP{2000MNRAS.318..163M}) is the closest to the TF analysis done here. In the absence of virial mass information the value of rotation velocity at a given luminosity does not suffice to make any general conclusions, so in general one has to make additional assumptions and/or modelling. For example, the stellar mass fraction in the halo $f_*$ obtained in previous work was lower, which lead to a higher virial mass for a given luminosity, which in turn requires lower concentrations and/or stellar mass to light ratios. By increasing $f_*$ close to its maximum value the virial mass can be reduced and this alleviates the problem. The same solution also solves the suggested overprediction of dark matter at the solar radius in our own galaxy \\cite{2001ApJ...554..114E}, since again if the stellar fraction is higher the virial mass can be lower (there may be additional problems for CDM profiles in the inner parts of our galaxy; e.g. \\citeNP{2001MNRAS.327L..27B}).  For early type galaxies it has been suggested that the concentrations are low from the strong lens statistics \\cite{2001ApJ...561...46K}, since very concentrated halos would overpredict the expected number of lenses. This is a difficult method to use since the expected number of lenses is very sensitive to the assumed luminosity function for early type galaxies as a function of redshift, which still has considerable uncertainty. Additional uncertainty arises from the adopted values for stellar mass to light ratios, which again can change the lensing statistics significantly. The lensing results are still compatible with low density CDM models, suggesting there is no discrepancy with our results, although more work is required to study this in detail.  There are other problems that have been suggested as troublesome for CDM, such as detailed shapes of velocity profiles in the optical region \\cite{2001ApJ...552L..23D}, bar rotation \\cite{2000ApJ...543..704D} or the halo structure of the Milky Way \\cite{2001MNRAS.327L..27B}. These probe inner regions of the galaxy where complicated physical processes may be taking place, so there is considerable more uncertainty in their theoretical predictions. For example, there are processes such as bar rotation that can disrupt dark matter cusps \\cite{2001astro.ph.10632W}. It has recently been shown that CDM profiles can fit most of the rotation curves for normal galaxies \\cite{2002astro.ph..1352J}. The galaxies that appear to be a problem for CDM belong to one of the specific subsamples, such as low surface brightness, dwarf or barred galaxies. It is possible, although not necessarily easy to arrange, that these samples are qualitatively different from the main population, for example by forming later and thus being less concentrated. Yet another possibility is that problems arise only below $L_*$, since our analysis is only valid for galaxies around and above $L_*$. Clearly, more work is required to resolve these issues. However, if the g-g lensing masses are correct, then for the main population of galaxies around and above $L_*$ the CDM model predictions for the amount of dark matter outside the inner few kpc do not exceed the observations, suggesting that the problems for CDM may not be as fundamental as previously suggested.  The author acknowledges the support of NASA, David and Lucille Packard Foundation and Alfred P. Sloan Foundation. I thank Mariangela Bernardi, \\v Zeljko Ivezi\\' c, Guinevere Kauffmann, Ravi Sheth and Tommaso Treu for useful discussions and Matthias Steinmetz for providing their manuscript prior to publication."
    },
    "0201/astro-ph0201499_arXiv.txt": {
        "abstract": "{ We present 850~$\\mu$m imaging polarimetry of the W51A massive star forming region performed with SCUBA on the JCMT. From the polarimetry we infer the column-averaged magnetic field direction, projected onto the plane of the sky. We find that the magnetic field geometry in the region is complicated. We compare the field geometry with 6~cm and CS~J=7-6 emission and determine that the magnetic field must be relatively weak and plays a passive role, allowing itself to be shaped by pressure forces and dynamics in the ionised and neutral gases. Comparisons are drawn between our data and 1.3~mm BIMA interferometric polarimetry data, from which we conclude that the magnetic field must increase in importance as we move to smaller scales and closer to sites of active star formation. ",
        "introduction": "The W51A cloud consists of a complex of {\\sc Hii} regions situated approximately 7-8 kpc away. It is one of the most active and luminous sites of massive star formation in the Galaxy with all the associated signatures: masers (Lepp\\\"anen et al \\cite{Leppanen98}; Zhang \\& Ho \\cite{Zhang95}; Genzel \\& Downes \\cite{Genzel77}), {\\sc Hii} regions (Scott \\cite{Scott78}), bright \\& dusty cores (Jaffe et al. \\cite{Jaffe84}; Genzel et al. \\cite{Genzel82}) and nearby supernova remnants (Koo \\& Moon \\cite{Koo97a,Koo97b}). Radio maps show a number of {\\sc Hii} regions, the most luminous being G49.5-0.4 (see e.g. Mehringer \\cite{Mehringer94}) which itself is dominated by the sources W51d and W51e. The former is directly associated with its infrared counterpart IRS2 (Goldader \\& Wynn-Williams \\cite{Goldader94}) and while W51e and the infrared source IRS1 are also closely associated, this section of the cloud is more complex.  A cluster of compact {\\sc Hii} regions are found to the east of W51e/IRS1, whose radio emission is dominated by the sources W51e2 and e1. The submillimetre emission is more closely associated with this cluster rather than with IRS1 and the peak of emission correlates with the position of W51e2 (Jaffe et al. \\cite{Jaffe84}), suggesting that W51e is a more evolved source while W51e1 and e2 are only now undergoing star formation. This is corroborated by the detection of infalling gas associated with W51e2 by Zhang \\& Ho (\\cite{Zhang97}).  Here, we present 850~$\\mu$m polarimetry of the W51A region revealing the magnetic field structure within it. There are a number of methods which allow for mapping the magnetic field geometry, and each has its own associated problems. By far the most promising method is through polarimetry of emission from dust grains which have been aligned (it is necessarily assumed) by interaction with the local magnetic field. Previous polarimetric observations of W51A show very low polarisation (Kane et al \\cite{Kane93}) consistent with a non-detection. However, these measurements were made with large beams ($\\sim$ 30\\arcsec) within which a significant degree of depolarisation can occur. For instance, Dotson et al. (\\cite{Dotson00}) present far-infrared polarimetry for W51A at 100 $\\mu$m taken with the KAO. Their results show low polarisation at the cores and higher degrees of polarisation in regions with more extended emission. Higher resolution and more sensitive measurements are clearly necessary to detect significant polarisation and to determine any structure in the polarisation pattern (e.g. see Lai et al. \\cite{Lai01} and Momose et al. \\cite{Momose01}). In their mid-infrared spectropolarimetric atlas of young stellar objects, Smith et al. (\\cite{Smith00}) measure 6\\% and 3\\% polarisation towards W51d in absorptive and emissive components, respectively, measurements made within 5\\arcsec~beams.  The importance of magnetic fields in star formation has become increasingly more apparent in recent years as computer models have become more sophisticated (e.g. Ostriker et al. \\cite{Ostriker01}). Magnetic fields are believed to play an important regulatory role during cloud collapse and accretion onto the protostar, as well as providing the driving and collimation mechanism for outflows. However, in order to comprehend the magnetic field's role properly, an empirical understanding of its behaviour is necessary.  In this paper, we present 850~$\\mu$m polarimetry not only of the dense cores in W51A but also of the surrounding more diffuse emission between these cores. In this fashion we investigate the relation between the magnetic field at the sites of star formation and the cold, expansive envelopes that surround them. By comparing this information with other studies we may also hope to understand what shapes the magnetic field. For instance, are the kinematics of gas and dust, as governed by gravitational collapse and the pressures within {\\sc Hii} regions, dominant, or is the energy density in the magnetic field of sufficient magnitude to control and regulate the dynamics? ",
        "conclusions": "We present 850~$\\mu$m polarimetry of the massive star forming region in W51A. The presented data has relatively high polarimetric signal-to-noise, which allows us to infer a magnetic field structure which is not uniform across our 2\\arcmin~field-of-view.  We have compared our data with a VLA 6~cm emission map as well as with a CS~J=7-6 map obtained from the JCMT archive. The comparison shows a correlation between the magnetic field geometry and the ionised and molecular gas. This indicates that the magnetic field is shaped by the gas dynamics and not the reverse, and therefore has a relatively weak magnetic field strength (parameterised by $\\beta~\\sim~1$).  However, Lai et al. (\\cite{Lai01}) have found from their 1.3~mm interferometer polarimetry that at smaller scales and towards the star forming cores the field lines are regular and ordered, implying that the magnetic field is comparatively strong ($\\beta~\\sim~0.01$) and plays a significant role in determining the gas dynamics -- although evidence suggests that further fragmentation and collapse probably occurred independent of the magnetic field. Nevertheless, when one looks at larger scales and further comparisons are made between the global field morphology and the gas dynamics, this picture breaks down.  Our data has shown how important it is to attempt to correlate magnetic field structures with the physical conditions of the (ionised and neutral) gas. In this way, one can develop a better understanding of the role magnetic fields play in star formation."
    },
    "0201/astro-ph0201516_arXiv.txt": {
        "abstract": "{We present spectroscopic measurements of the 23 new candidate members of the TW~Hydrae Association from \\citet{MF}. Based on H$\\alpha $ and Li~6708~\\AA{} strengths together with location on a color-magnitude diagram for Hipparcos TWA candidates, we found only three possible new members (TYC~7760-0835-1, TYC~8238-1462-1, and TYC~8234-2856-1) in addition to the already known member, TWA~19. This eliminated most of the candidates more distant than 100~pc. Three Tycho stars, almost certainly members of the Lower Centaurus Crux association, are the most distant members of the TWA.  A claim of isotropic expansion of TWA has to be re-evaluated based on our new results. Generally, one cannot identify new members of a diffuse nearby stellar group based solely on kinematic data. To eliminate interlopers with similar kinematics, spectroscopic  verification is essential. ",
        "introduction": "During the past few years, a handful of young ($<100$~Myrs) and nearby ($\\leq 100$~pc) groups of stars have been identified. Usually, they occupy a few hundred square degrees in the plane of the sky (thus are less conspicuous than compact clusters) and are not associated with any known interstellar clouds. Therefore they have been hard to find and have remained unnoticed until recently. Beginning with the TW Hydrae Association \\citep[\"TWA\", ][and references therein]{TWA}, currently, there are about eight such associations reported in the literature. For the most up-to-date information on these groups of stars, readers are refered to \\citet{YStars}.  Thermal emission from massive planets can be detectable when they are separated far enough from the primary stars. For a given linear separation between primary and planet, a closer system to Earth will provide a larger angular separation making detection of the planet easier. Planets around young stars are still hot and they emit a fair amount of infrared radiation. Therefore, young and nearby stars are excellent targets for observational searches for giant planets and dusty ``proto-planetary'' disks.  The TWA (d$\\sim 60$~pc), the $\\beta $~Pictoris moving group (d$\\sim 35$~pc), and the Tucana/HorA Association (d$\\sim 45$~pc) are the three youngest and closest groups to Earth; hence they form the best targets for ground- and space-based searches for planets. The Capricornus group \\citep{vandenAncker} is also a very young nearby group, however it is almost certainly a subgroup of the bigger $\\beta $~Pictoris group \\citep{bPic2}.  There are currently 19 known star systems in TWA and most of its members have been observed with the  Hubble Space Telescope and/or with a ground-based telescope with adaptive optics (AO) capability. Some TWA members possess substellar companions (TWA~5; \\citealt{TWA5}) or prominent dusty disks (TW~Hydrae, HR~4796A, Hen3-600, and HD~98800). Any newly identified members of TWA, therefore, are excellent targets for planet or planetary disk searches.   \\def\\mmc{\\multicolumn{2}{c}} \\begin{table*}  \\caption{Spectroscopic data for Makarov and Fabricious TWA candidates.}  {\\centering \\begin{tabular}{clcr@{$\\pm$}lr@{$\\pm$}lr@{$\\pm$}lcccc} \\hline  Name&\\multicolumn{1}{c}{ $(B-V)^{\\dag}$}& R.V.& \\multicolumn{6}{c}{R.V. measured (km/sec)}& EW(H$\\alpha $)& EW(Li)& $v\\sin i$& TWA\\\\ \\cline{4-9} HIP or TYC&\\multicolumn{1}{c}{(error)}& predicted& \\mmc{SSO}& \\mmc{Torres et al.}& \\mmc{Others$^{\\ddag }$}& (\\AA{})& (m\\AA{})& (km/sec)& member?\\\\ \\hline  53911      & 0.64(12)    & 12.7& \\mmc{--}    & $+$12.9&0.3 & \\mmc{--}    & $-220$  & 390  & 4  & TWA 1  \\\\ 7201-0027-1& 1.88(37)    & 10.6& \\mmc{--}    & $+$11.2&0.3 & \\mmc{--}    & $-1.89$ & 490  & 13 & TWA 2  \\\\ 55505      & 1.19(03)    &  9.1& \\mmc{--}    &  $+$9.2&0.2 & \\mmc{--}    & $0$     & 360  & -- & TWA 4  \\\\ 7223-0275-1& 1.44(33)    & 10.0& $-$30.6&6.6 & $+$6.9:&2.0 & \\mmc{--}    & $-13.4$ & 570  & 36:& TWA 5  \\\\ 7183-1477-1& 1.61(26)    & 16.2& \\mmc{--}    & \\mmc{--}    & \\mmc{--}    & $-4.65$ & 560  & -- & TWA 6  \\\\ 7190-2111-1& 1.31(18)    & 11.0& \\mmc{--}    & \\mmc{--}    & \\mmc{--}    & $-4.95$ & 440  & -- & TWA 7  \\\\ 57589      & 1.66(40)$^G$& 12.6& \\mmc{--}    & $+$10.2&0.4 & \\mmc{--}    & $-5.01$ & 480  & -- & TWA 9  \\\\ 61498      & 0.00(01)$^G$& 10.5& \\mmc{--}    & \\mmc{--}    &  $+$9.4&2.3 & --      & --   & -- & TWA 11 \\\\ 57524      & 0.60(02)    & 20.0& $+$11.5&3.8 & \\mmc{--}    & \\mmc{--}    & $0.57$  & 189  & 24 & TWA 19 \\\\ \\hline 46535      & 0.50(01)$^G$& 16.0& $+$18.0&7.0 & \\mmc{--}    & $+$22.1&5.0 & $1.43$  & 63   & 37 & no     \\\\ 47039      & 0.42(01)$^G$& 16.9& $+$10.9&1.3 & \\mmc{--}    & $+$12.2&0.4 & $1.63$  & $<10$&  8 & no     \\\\ 48273      & 0.48(01)$^G$& 10.7& $-$31.4&2.9 & $+$16.2&0.1 & \\mmc{$+17.0^{\\S }$}&2.0& 21 \\& 22 & 9 \\& 14& no\\\\ 0829-0845-1& 0.69(05)    & 20.4&  $+$9.5&6.5 & \\mmc{--}    & \\mmc{--}    & $0.85$  & $<10$& 30 & no     \\\\ 6604-0118-1& 1.07(06)    & 16.5&\\mmc{$-19,+54,+69$}& $+$27.0&0.3 & \\mmc{--}& $-0.5$& 93   & 20 & no     \\\\ 49530      & 0.94(01)$^G$& 20.7& $+$22.5&1.1 & \\mmc{--}    & $+$16.7&2.0 & $1.17$  & 34   & 2  & no     \\\\ 6625-1087-1& 0.87(13)    & 21.3& $+$16.1&1.5 & \\mmc{--}    & \\mmc{--}    & $0.28$  & 141  & 10 & no     \\\\ 7178-1493-1& 0.73(18)    & 15.2& $+$48.0&3.0 & \\mmc{--}    & \\mmc{--}    & $0.17$  & $<10$& 15 & no     \\\\ 7183-1879-1& 0.96(20)    & 18.9& $+$11.0&1.6 & \\mmc{--}    & \\mmc{--}    & $0.74$  & 79   & 11 & no     \\\\ 7188-0575-1& 1.11(05)    & 12.2& $+$42.9&2.1 & \\mmc{--}    & \\mmc{--}    & $-0.28$ & 50   & 25 & no     \\\\ 50796      & 1.19(02)$^G$& 13.1& $+$22.4&0.9 & $+$13.1&1.0 & \\mmc{--}    & $0.20$  & $<10$& 8  & no     \\\\ 7710-2231-1& 1.07(03)    & 22.5& $-$12.2&1.0 & \\mmc{--}    & \\mmc{--}    & $0.63$  & $<14$& 7  & no     \\\\ 52462      & 0.87(01)$^G$&  9.0& $+$22.7&0.5 & \\mmc{--}    & \\mmc{--}    & $0.78$  & 150  & 1  & no     \\\\ 52787      & 0.82(02)    &  9.9& $+$24.0&0.6 & \\mmc{--}    & \\mmc{--}    & $0.82$  & $120$& 2  & no     \\\\ 53486      & 0.91(02)$^G$&  3.7& $+$4.3 &1.0 &  $+$5.5&0.3 & \\mmc{--}    & $0.88$  & $<10$& 1  & no     \\\\ 55899      & 0.07(02)    & 21.8& \\mmc{--}    & \\mmc{--}    & \\mmc{--}    & $2.0$   & --   & 210& no?    \\\\ 57129      & 0.54(03)    & 12.9& $+$15&20    & \\mmc{--}    & $+$25.6&10  & $1.56$  & $<10$& 190& no     \\\\ 57269      & 0.91(01)$^G$&  8.5& \\mmc{--}    & \\mmc{--}    & \\mmc{$+15.9^{\\star }$}&--& 196& 20& no     \\\\ 59315      & 0.71(01)$^G$&  6.6& $+$15.9&0.6 & \\mmc{--}    & \\mmc{--}    & $0.83$  & 151  & 5  & no     \\\\ 7760-0835-1& 0.51(03)    & 19.4& $+$10.0&2.6 & \\mmc{--}    & \\mmc{--}    & $0.83$  & 161  & 12 & yes?   \\\\ 8238-1462-1& 0.77(04)    & 16.8& $+$12.0&3.0 & \\mmc{--}    & \\mmc{--}    & $0.36$  & 294  & 18 & yes?   \\\\ 8234-2856-1& 0.81(05)    & 20.1& $+$13.2&2.4 & \\mmc{--}    & \\mmc{--}    & $-0.49$ & 342  & 16 & yes?   \\\\ \\hline \\multicolumn{13}{l}{$^{\\dag }$ estimated from Tycho $(B_{T}-V_{T})$ using a relation given in \\citet{Bessell}.}  \\\\ \\multicolumn{13}{l}{~~ Stars with superscript \"$G$\", $(B-V)$ colors are from ground photometry adapted from Hipparcos Input Catalog.}\\\\ \\multicolumn{13}{l}{$^{\\ddag }$ radial velocities from \\citet{B-B} unless noted otherwise.}\\\\ \\multicolumn{13}{l}{$^{\\S }$ radial velocity from \\citet{HIP48273}.}\\\\ \\multicolumn{13}{l}{$\\star $ radial velocity from \\citet{HIP57269}.}\\\\ \\end{tabular}\\par} \\end{table*}  Using a modified convergent point method, \\citet[MF hereafter]{MF} found 23 new kinematic candidate members of TWA beside eight known members. In fact, one of the 23 candidates is a known TWA member, TWA~19 \\citep[HIP\\,57524;][]{TWA}. Based on a kinematic model of TWA including the candidate members, MF predicted radial velocities of the 23 new candidates. They also suggest that TWA expands with a rate of $0.12\\, \\mathrm{km}/\\mathrm{sec}\\cdot \\mathrm{pc}^{-1}$ and that TWA is part of a larger structure, the ``Gould disk'', rather than a separate group of stars. Considering the importance of TWA members as aforementioned observational targets and as corner-stones in investigations of the surrounding environment (local association, Sco--Cen complex, etc.), all of the MF candidates need to be verified spectroscopically. In this paper, we present spectroscopic data for all MF candidates and discern true members from interlopers based on spectroscopic youth indicators. ",
        "conclusions": ""
    },
    "0201/astro-ph0201270_arXiv.txt": {
        "abstract": "We present the distribution of a statistical sample of nearby galaxies in the $\\kappa$-space ($\\kappa_1 \\propto log~M$, $\\kappa_2 \\propto log~I_e$, $\\kappa_3 \\propto log~M/L$). Our study is based on near-IR (H-band: $\\rm \\lambda = 1.65~\\mu m$) observations, for the first time comprising early- and late-type systems. Our data confirm that the mean effective dynamical mass-to-light ratio $M/L$ of the E+S0+S0a galaxies increases with increasing effective dynamical mass $M$, as expected from the existence of the Fundamental Plane relation. Conversely, spiral and Im/BCD galaxies show a broad distribution in $M/L$ with no detected trend of $M/L$ with $M$, the former galaxies having $M/L$ values about twice larger than the latter, on average. For all the late-type galaxies, the $M/L$ increases with decreasing effective surface intensity $I_e$, consistent with the existence of the Tully--Fisher relation. These results are discussed on the basis of the assumptions behind the construction of the $\\kappa$-space and their limitations. Our study is complementary to a previous investigation in the optical (B-band: $\\rm \\lambda = 0.44~\\mu m$) and allows us to study wavelength-dependences of the galaxy distribution in the $\\kappa$-space. As a first result, we find that the galaxy distribution in the $\\kappa_1$--$\\kappa_2$ plane reproduces the transition from bulge-less to bulge-dominated systems in galaxies of increasing dynamical mass. Conversely, it appears that the $M/L$ of late-types is higher (lower) than that of early-types with the same $M$ in the near-IR (optical). The origins of this behaviour are discussed in terms of dust attenuation and star formation history. ",
        "introduction": "The study of the scaling relations between photometric and kinematical properties of present-day stellar systems is a powerful tool for the comprehension of the processes that led to their formation and evolution (e.g. Mao \\& Mo 1998; Mao, Mo \\& White 1998; Pahre, de Carvalho \\& Djorgovski 1998; Avila-Reese \\& Firmani 2000). The relation between the total luminosity $L$ and the central velocity dispersion $\\sigma_0$ (Faber \\& Jackson 1976) and the relation between the effective radius $r_e$ (the radius that contains half of the galaxy total luminosity) and the central surface brightness $\\mu_0$ (Kormendy 1977) represent benchmarks for any theoretical modelling of the early-type galaxies (de Zeeuw \\& Franx 1991 and references therein). Further analyses pointed out that these galaxies and the bulges of the spiral galaxies populate a two-dimensional manifold in the parameter space defined by $r_e$, $\\mu_e$ and $\\sigma_0$, named the Fundamental Plane (FP) (Djorgovski \\& Davis 1987; Dressler et al. 1987), where $\\mu_e$ is the mean effective surface brightness (mean surface brightness within $r_e$). The profound implications of this result, its interpretation and its extension to parent systems at high redshifts are still grounds both of observational challange and of intellectually exciting debate (e.g., Scodeggio et al. 1998; Pahre, Djorgovski \\& de Carvalho 1999 and references therein). For the late-type galaxies, the correlation between the total luminosity $L$ and the maximum rotational velocity $V_{Max}$ (Tully \\& Fisher 1977), as well as the distribution in the parameter space defined by exponential disk scale-length $h$ and central surface brightness $\\mu_0$ (Grosb$\\rm \\o$l 1985; de Jong 1996a; see also Pierini 1997) are now compelling for any theory of disk-galaxy formation (e.g. Dalcanton, Spergel \\& Summers 1997; Mo, Mao \\& White 1998; Syer, Mao \\& Mo 1999). In order to reproduce the locus of these systems in the $h$--$\\mu_0$ plane, a bivariate distribution in mass and spin parameter of the proto-galaxies is required, whatever the model of disk-galaxy formation is (cf. Dalcanton, Spergel \\& Summers 1997; Firmani \\& Avila-Reese 2000).  The effective and central photometric parameters involved by the previous scaling relations were determined traditionally from fitting either a de Vaucouleurs $r^{1/4}$ law (de Vaucouleurs 1948) or an exponential law (Freeman 1970) to the radial surface brightness profiles. The first law was used to describe the light distribution along the radial coordinate of the elliptical galaxies and of the bulges of the spiral galaxies, after the pioneering analyses of de Vaucouleurs (1948, 1959). The second one was found to be a good description of the light profiles of the disks, also in the case of S0 galaxies (Burstein 1979).  Burstein et al. (1997 -- hereafter referred to as BBFN; see also Bender, Burstein \\& Faber 1992; Burstein et al. 1995) proposed a three-dimensional parameter system, called $\\kappa$-space, in order to represent and to compare the global relationships of any stellar system (i.e., from the Galactic globular clusters to the galaxy clusters), in a self-consistent way. The axes of the $\\kappa$-space, $\\kappa_1$, $\\kappa_2$ and $\\kappa_3$, are proportional to the logarithms of the galaxy mass and mass-to-light ratio and of a third quantity, which is basically the surface brightness, respectively, and do not depend on the fitting law of the surface brightness profile. They reflect a set of physical assumptions on the structure and velocity pattern of the different self-gravitating, equilibrium stellar systems (see BBFN), which may be subject of criticism on the basis of some observational results (e.g. Caon, Capaccioli \\& D'Onofrio 1993; Busarello et al. 1997). BBFN investigated kinematical and structural properties of nearby galaxies on the basis of a large optical (B-band: $\\rm \\lambda = 0.44~\\mu m$) data-set. Further analyses focussed on the distribution in the near-IR $\\kappa$-space either of present-day E+S0 systems (Pahre, Djorgovski \\& de Carvalho 1998 -- hereafter referred to as PDdC) or of disks and bulges of present-day late-type galaxies (Moriondo, Giovanelli \\& Haynes 1999). The adoption of near-IR photometry gives an insight of the scaling properties of galaxies less biased by differences in age and metallicity of the stellar populations (de Jong 1996b) and by dustiness (Cardelli, Clayton \\& Mathis 1989 and references therein) than the picture emerging from optical studies.  Here we present the near-IR (H-band: $\\rm \\lambda =1.65~\\mu m$) $\\kappa$-space of both early- and late-type galaxies as a whole at z=0. In Sect. 2 we define the $\\kappa$-space axes and their relations with typical galaxy parameters, such as mass, radius, surface brightness and luminosity, and discuss the physical assumptions behind the $\\kappa$-space and their limitations. The selection criteria of our sample and the determination of the galaxy parameters involved by the definition of the $\\kappa$-space axes are described in Sect. 3. Here we note that this sample comprises several Virgo cluster galaxies down to $\\rm M_B = -15~mag$ (cf. Boselli et al. 1997, 2000a), though it is not complete to low surface brightness and dwarf galaxies. We present and discuss the galaxy distribution in the near-IR $\\kappa$-space in Sect. 4. Discussion and conclusions on the wavelength-dependence of the galaxy distribution in the $\\kappa$-space are given in Sect. 5. ",
        "conclusions": "The existence of the Hubble sequence of galaxies (Hubble 1926; Sandage 1961) reflects differences in structure, kinematics and global star-formation histories of the classified stellar systems (Roberts \\& Haynes 1994). The study of the galaxy distribution in the so-called $\\kappa$-space (BBFN) contributes to the understanding of the galaxy phenomenology, under the non trivial assumption that early- and late-type galaxies form two distinct homologous classes in gravitational equilibrium. Our investigation of the $\\kappa$-space of nearby galaxies is based on near-IR surface photometry and is complementary to the near-IR studies of PDdC (limited to ellipticals and lenticulars) and Moriondo, Giovanelli \\& Haynes (1999) (limited to bulges and disks of late-type galaxies) and to the optical study of BBFN, though the photometric parameters are derived via different procedures in these four studies.  As a first result, we find that galaxies of the same morphological type distribute in the $\\kappa_1$--$\\kappa_2$ projection of the $\\kappa$-space in a way analogous to the optical case. This distribution reproduces the ``generalized'' Kormendy relation (BBFN) since it finds analogies with the Kormendy relation of elliptical and lenticular galaxies (Kormendy 1977, 1988; Burstein 1979) on one hand and, on the other, the distribution of the disk-components of late-type galaxies in the $h$--$\\mu_0$ plane (Grosb$\\rm \\o$l 1985; de Jong 1996a; Pierini 1997). It traces changes in the global structure associated with the dynamical mass of the galaxy (cf. Paper V and IX), beyond any possible systematic effect due to differences in the radial distribution of the stellar populations which dominate the light emission in different pass-bands, whether these differences are intrinsic or due to differential attenuation by dust (cf. Witt, Thronson \\& Capuano 1992 for the early-type galaxies). We also confirm that the distribution of the early-type galaxies in the $\\kappa_1$--$\\kappa_3$ plane reproduces their Fundamental Plane relation and that, conversely, the distribution of the late-type galaxies in the $\\kappa_2$--$\\kappa_3$ plane reproduces the Tully--Fisher relation of the latter stellar systems. This time, the galaxy distribution in the $\\kappa_1$--$\\kappa_3$ and $\\kappa_2$--$\\kappa_3$ projections of the $\\kappa$-space depends on the pass-band adopted in order to determine the photometric properties of the galaxies, as a consequence of the dominant role of the mass-to-light ratio.  The Fundamental Plane relation (in $\\kappa$-space notation) shows that the dynamical mass-to-light ratio of elliptical and lenticular galaxies increases with dynamical mass, both in the optical and in the near-IR. The neglect of dust effects, even for rather modest amounts of dust, leads to an overestimate of the total dynamical mass-to-light ratio by a factor of $\\sim$ 20\\% per optical depth unit (Baes, Dejonghe \\& Rijcke 2000). Therefore, this relation holds almost unaffected by dust bias of the photometric properties in the near-IR. By contrast, the neglect of the rotational velocity on the mass estimate affects the Fundamental Plane relation, whatever the photometric pass-band is. We estimate that the systematic underestimate of the dynamical mass of the ellipticals is about 10\\%, on average. As a consequence, in the near-IR, a giant elliptical galaxy would have a corrected mean effective mass-to-light ratio of about 1.6 (in solar units), on average, still lower than the average mean effective mass-to-light ratio of late-type galaxies. The correction for the lenticular galaxies is not straightforward, but we may expect that it ranges between 10 and 30\\%, on average, the latter value applying to rotating ellipticals of maximum ellipticity $\\sim$ 0.4. As a consequence, in the near-IR, the corrected mean effective mass-to-light ratios of lenticular galaxies might be intermediate between those of earlier and later Hubble types.  Spiral and Im/BCD galaxies show very different behaviours in the optical and near-IR $\\kappa_1$--$\\kappa_3$ planes. First of all, they have mean effective mass-to-light ratios higher than those of elliptical and lenticular galaxies of the same mass in the near-IR but it is vice versa in the optical (cf. BBFN). A study of the non-trivial effects of dust attenuation on the surface photometry of the late-type galaxies is beyond the reach of this analysis and, therefore, we have assumed statistically-based inclination corrections of the near-IR photometric parameters of these stellar systems (Equ. 10 and 11). In the near-IR, these corrections produce an increase of the mean effective mass-to-light ratio (Fig. 3a,b), opposite to what expected for early-type galaxies, so that we are confident that the relative distribution of early- and late-type galaxies $\\kappa_1$--$\\kappa_3$ plane is robust. Second, in the optical, the mass-to-light ratio vs. mass relation shows a strong dependence on the Hubble type of the spiral and Im/BCD galaxies. In the near-IR this dependence is not detected (if any) though there may be a hint that Im/BCD galaxies have higher mass-to-light ratios than spirals.  If not due to systematic differences either in the statistically-based corrections for dust effects and inclination adopted by us and BBFN or in the estimate of the photometric parameters, the different locations of early- and late-type galaxies in the optical and near-IR $\\kappa_1$--$\\kappa_3$ planes may originate from residual dust effects and/or from differences in the characteristic stellar populations, both dependent on morphology and mass. The stellar mass-to-light ratio of a simple stellar population (SSP) increases with age both in the optical and in the near-IR, its early-time evolution being very fast until $\\sim 5$ Gyrs and becoming mild afterwards (e.g. Maraston 1998). In particular, the stellar mass-to-near-IR light ratio reaches a value within 20\\% from the present one ($\\sim 1.2$ at 15 Gyrs) after only 3 Gyrs, while the stellar mass-to-optical light ratio amounts to 80\\% of the present one ($\\sim 10$) after 13 Gyrs. Of course a galaxy is not reproduced by a SSP but by a mix of different SSPs, weighted by its global star formation history. In particular, star formation activity leads to dust production, so that the mass-to-light ratio of a more recently born SSP will be reduced by dust attenuation both at the stellar photosphere and in the interstellar medium. It is commonly accepted that the bulk stellar population of giant elliptical and lenticular galaxies is older than those of spiral and Im/BCD galaxies, i.e. that the former have transformed gas into stars much faster than the latter (e.g. Renzini 1998). An analogous trend of decreasing ages of the characteristic stellar population with later Hubble type is suggested for the late-type galaxies (e.g. Sandage 1986; Kennicutt, Tamblyn \\& Congdon 1994; Gavazzi \\& Scodeggio 1996; Boselli et al. 2001). If these considerations apply, we expect that the early-type galaxies have higher stellar mass-to-light ratios than later ones, where star formation is still going on, under the assumptions that galaxies of all morphological types have the same age and the same initial mass function. In particular, the stellar mass-to-near-IR light ratios of nearby early- and late-type galaxies will not differ much if the peak of star formation activity took place more than 3 Gyrs ago for all of them. By contrast, the stellar mass-to-optical light ratios of early- and late-type galaxies may still differ by a maximum factor of $\\sim$ 7 (Maraston 1998), if this peak took place not less than 3 Gyrs ago for all of them.  We observe that the dynamical mass-to-near-IR light ratio of elliptical and lenticular galaxies is lower than that of spiral of the same dynamical mass. From the previous considerations, we conclude that the dynamical-to-stellar mass ratio of the former galaxies is lower than that of the latter, if the peak of the star formation activity took place more than 3 Gyrs ago for all of them. Extending the same conclusion to Im and BCD galaxies is dangerous, since the near-IR luminosity may be seriously contaminated by emission from younger stellar populations. Since BBFN observe that the dynamical mass-to-optical light ratio of elliptical and lenticular galaxies is higher than that of spiral galaxies of the same dynamical mass, this behaviour is consistent with the previous conclusion if the difference in the stellar mass-to-optical light ratio, due to the different star formation histories of these two classes of stellar systems, is larger than the difference in the dynamical-to-stellar mass ratio.  Finally, we note that the optical and near-IR Fundamental Plane relations should be analogous if the peak of star formation took place more than 3 Gyrs ago for all the elliptical and lenticular galaxies, but for a color term, when the differential effects of diffuse dust on the photometric parameters are taken into account. Whether this color term is not or is constant, it depends on the existence of the color-magnitude relation or not (Scodeggio 2001 and references therein). The answer to this question would help understanding whether the Fundamental Plane relation (in $\\kappa$-space notation) is (mainly) due to the increase either of the dynamical-to-stellar mass ratio or of the stellar mass-to-light ratio with dynamical mass."
    },
    "0201/astro-ph0201046_arXiv.txt": {
        "abstract": "We derive similarity solutions for the expansion of negative initial density perturbations $\\delta M/M \\propto r^{-\\sp}$ $(\\sp>0)$ in an Einstein-de Sitter universe filled with collisional baryonic gas and collisionless matter.  The solutions are obtained for planar, cylindrical, and spherical perturbations.  For steep perturbations, central cavities surrounded by over-dense shells expanding as $t^{{2(s+1)}/{(3s)}}$ develop in both matter components.  We also consider the case when baryonic shells are driven by internal pressure.  Before redshift $\\sim 8-10$ baryonic shells cool by inverse Compton scattering with the cosmic background radiation and can form bound objects of mass $\\sim 10^{11}\\Omega_{\\rm b}^{-1/2} M_{\\odot} (1+z)^{-3/2}$.  Hot shocks of radius $R$ cause a Compton y-distortion of order $y=10^{-5}\\Omega_{\\rm b} (R/10 {\\rm Mpc})^2 f_F$, where $f_F$ is the volume filling factor of these shells. ",
        "introduction": "\\label{equations}  We write the Newtonian equations of motion governing the adiabatic evolution of symmetric perturbations in a collisional fluid (gas) of adiabatic index $\\gamma$. We assume that the expansion scale factor of the universe is $a(t)\\propto t^{2/3}$, where $t$ is the age of the Universe, the Hubble function is $H(t)=2/(3t)$, the total background density is $\\rho_c=3H^2/(8\\pi G)=1/(6\\pi G t^2)$, and the ratio of the mean collisional baryonic density to the critical density $\\rho_c$ is denoted by $\\Omega_{\\rm b}$. Although self-similar evolution exists for any $\\Omega_{\\rm b}\\le 1$, in this paper we restrict the analysis to either $\\Omega_{\\rm b}=1$, or $\\Omega_{\\rm b}\\ll 1$.   Denote by $r$ and $\\upsilon\\equiv \\dd r/ \\dd t$ the physical position and velocity of a gas shell, where $ r=0$ is the symmetry center of the perturbation.  Further, let $\\rho(r,t)$ and $p(r,t)$ be the gas density and pressure at $r$.  We write the mass within a distance $r$ from the symmetry center as $m(r,t)=\\int_0^r x^{\\nn-1}\\rho(x,t)\\dd x$, where $\\nn=1,2$, and 3 refer, respectively, to planar, cylindrical, and spherical perturbations.  The mass within a fixed shell varies with time like $m \\sim t^{-2(3-\\nn)/3}$, because of the Hubble expansion along $3-\\nn$ of the axes.  At this stage we neglect such physical effects as viscosity, thermal conduction, cooling etc., so the equations of motion of collisional fluid are, the continuity equation,  \\def\\ff{\\frac{2(3-\\nn)}{3}} \\begin{equation} \\frac{\\dd\\left[\\rho t^{\\ff}\\right]}{\\dd t}=-t^\\ff \\rho r^{1-\\nn}\\partial_r(r^{\\nn-1}\\upsilon) \\; , \\label{eom1} \\end{equation} the  Euler equation, \\begin{equation} \\frac{\\dd \\upsilon}{\\dd t}-\\frac{2}{9}\\frac{3-\\nn}{\\nn}\\frac{r}{t^2} =-\\frac{\\partial_r p}{\\rho}- \\frac{4\\pi G m_x}{r^{n-1}} \\; , \\label{eom2} \\end{equation} the  adiabatic condition,  \\begin{equation} \\frac{\\dd}{\\dd t}(p\\rho^{-\\gamma})=0 \\; , \\label{eom3} \\end{equation} and the relation, \\begin{equation} \\partial_r m= r^{\\nn -1}\\rho \\; . \\label{eom4} \\end{equation}  In equation (\\ref{eom2}), $m_x$ stands for total (collisional and collisionless) mass.  There are several similar techniques, which we do not describe here in detail, for calculating the collisionless mass profile (Bertschinger 1985; Fillmore \\& Goldreich 1984). The main idea is that, since the motion is self-similar, knowing the orbit of a single particle one can calculate the mass profile and vice versa. Given the initial guess for mass profile, after iterative calculation of mass profile and orbit one arrives to the self-similar solution.  The initial conditions leading to self-similar expansion are specified at an early time close to zero, $t_i$, as \\begin{eqnarray} \\label{inid} \\frac{\\delta M}{M}&=&-\\left(\\frac{r}{r_0}\\right)^{-\\sp} \\; ,r\\gg r_0 \\\\ \\label{inivel} \\upsilon(r,t_i)&=&\\frac{2}{3t_i}r \\; , \\\\ \\label{inip} p(r,t_i)&=&0 \\; , \\label{inic} \\end{eqnarray} where $s>0$ and $\\delta M/M$ is the mean density contrast interior to $r$. The Einstein-De Sitter Universe and the initial conditions are completely scale-free.  The only characteristic length in the evolution of the perturbation is the scale of non-linearity $r_*(t)$. For self-similar collapse of positive perturbations the turn-around radius is usually used as the scale of non-linearity. For a negative perturbation, $r_*$ can be defined at any time as the radius interior to which the mean density contrast has a certain fixed value.  This means that \\begin{equation} r_* \\propto t^\\alpha \\qquad , \\qquad \\alpha=\\frac{2s}{3(s+1)}\\; . \\label{ralpha} \\end{equation} This is also the way the turnaround radius in positive perturbation depends on time. The choice of proportionally factor in (\\ref{ralpha}) is arbitrary. Here it is chosen such that a particle with initial radius $r_i$ reach $r_*$ at $(6\\delta M/M)^{-3/2}t_i$, where the factor of 6 is arbitrary. For this choice of $r_*$ the mean density contrast interior to $r_*$ is approximately $-0.09$ in all symmetries.  A consequence of  self-similarity is  that the partial differential equations of motion can be transformed into ordinary differential equations. This is done by working with  $\\lambda\\equiv r/r_*$ and the dimensionless fluid variables,  \\begin{eqnarray} \\label{scalev}\\upsilon(r,t)&=&\\frac{r_*}{t}V(\\lambda)\\\\ \\label{scaled}\\rho(r,t)&=&\\Omega_{\\rm b}\\rho_c D(\\lambda)\\\\ \\label{scalep}p(r,t)&=&\\Omega_{\\rm b}\\rho_c\\left(\\frac{r_*}{t}\\right)^2 P(\\lambda)\\\\ \\label{scalem}m(r,t)&=&\\frac{1}{3}\\Omega_{\\rm b}\\rho_c r_*^\\nn M(\\lambda) \\; . \\end{eqnarray}  Expressed in terms of these variables, the equations (\\ref{eom1}-\\ref{eom4}) become, respectively, \\begin{equation} \\label{a1} \\left(V-\\alpha\\lambda\\right)D'+\\left(\\frac{\\nn-1}{\\lambda}V+V'- \\frac {2\\nn}{3}\\right)D=0 \\; , \\end{equation} \\begin{equation} \\label{a2}\\left(\\alpha-1\\right)V+\\left(V-\\alpha\\lambda\\right)V' -\\frac{2}{9}\\frac{3-\\nn}{\\nn}\\lambda = -\\frac{P'}{D}-\\frac{2}{9}\\frac{M}{\\lambda^{\\nn-1}} \\; , \\end{equation}  \\begin{equation} \\label{a3}\\left(\\gamma\\frac{D'}{D}- \\frac{P'}{P}\\right)\\left(V-\\alpha\\lambda\\right)=2\\left(\\alpha-2+\\gamma\\right) \\; , \\end{equation}  \\begin{equation} \\label{a4}M'=3\\lambda^{\\nn-1}D \\; , \\end{equation} where the prime symbol denotes derivatives with respect to $\\lambda$.  Self-similarity implies that the shock appears (if it does) at fixed $\\lambda=\\lambda_{\\rm s}=r_{\\rm s}/r_*$, so the physical radius of the shock $r_{\\rm s}\\propto t^\\alpha$ and its non-dimensional speed is $(r_*/t)^{-1}(\\dd r_{\\rm s}/\\dd t)= \\alpha\\lambda_{\\rm s}$.  At the surface of the shock the fluid variables satisfy the jump conditions obtained from mass, momentum, and energy conservation (cf. Spitzer 1978, \\S 10.2):  \\begin{eqnarray} \\frac{\\upsilon'_+}{\\upsilon'_-}&=&\\frac{\\rho_-}{\\rho_+}=\\frac{\\gamma-1}{\\gamma+1},\\\\ p_++\\rho_+\\upsilon_+^{'2}&=&p_-+\\rho_-\\upsilon_-^{'2},\\\\ \\upsilon_-&=&\\upsilon'_-+\\frac{dr_{\\rm s}}{dt},\\\\ \\upsilon_+&=&\\upsilon'_++\\frac{dr_{\\rm s}}{dt}, \\end{eqnarray} where the superscripts of the minus and plus signs refer to pre- and post-shock quantities and $\\upsilon'$ is the velocity relatively to the shock position.  In terms of the non-dimensional fluid variables we obtain \\begin{eqnarray} \\label{jump1} V^+&=&\\alpha\\lambda_{\\rm s}+\\frac{\\gamma-1}{\\gamma+1}(V^- -\\alpha\\lambda_{\\rm s})\\; ,\\\\ \\label{jump2}D^+&=&\\frac{\\gamma+1}{\\gamma-1}D^- \\; ,\\\\ \\label{jump3}P^+&=&\\frac{2}{\\gamma+1}D^-(V^--\\alpha\\lambda_{\\rm s})^2\\; . \\end{eqnarray} Outside the shock, at $\\lambda >\\lambda_{\\rm s}$, the pressure vanishes and the pre-shock fluid variables can be found by solving the equations (\\ref{a1}-\\ref{a4}) with zero pressure. Analytic solutions for zero pressure exist for planar and spherical geometries, but not for cylindrical (Zel'dovich 1970, Peebles 1980, Bertschinger 1985).  \\subsection{The existence of shocks and the value of $s$} Lower and upper limits on $s$ exist for a shocked expansion of an isolated negative cosmological perturbation.  A necessary condition for the formation of a shock is that the perturbation is steep enough so that inner shells expand faster than outer shells.  An inspection of the equations of motion shows that this happens only if $(\\alpha-1)\\alpha< \\frac{2}{9}\\frac{3-n}{n}$, that is $s>2$ for spherical, $s>1.54$ for cylindrical, and $s>1$ for planar symmetry. This lower limit disappears if an external energy source is placed at the center of the perturbation.  An upper limit on $s$ also exists. To see this, consider the case of a spherical perturbation.  If the perturbation is negative, the total energy of a particle is positive and so the energy inside the shock must increase as it sweeps more particles.  On the other hand, self-similarity implies that the total energy within $\\lambda_{\\rm s}$ varies with time as $t^{5\\alpha-4}$ which is a decreasing function of time for $\\alpha < 4/5$, i.e., $s< 5$. To satisfy both conditions the energy within $\\lambda_{\\rm s}$ must be negative. The only way to construct such system is by placing a bound object (i.e., negative energy) at the center. However, since there is no way to make this object grow with time, the evolution of the perturbation never becomes self-similar. The limiting case of $s\\rightarrow 5$, which corresponds to a  shock in a homogeneous medium (compensated void), was investigated by Bertschinger (1983). The corresponding upper limits for cylindrical and planar perturbations are  $s\\approx 3.5$ and $s=2$, respectively.   \\subsection{Boundary conditions} Solving the self-similar equations (\\ref{a1}--\\ref{a4}) requires boundary conditions on the fluid variables.  We will see in the next section that shocked expansion is associated with a dense shell whose outer boundary is the shock. Outside the shock the pressure is zero and the equations can be solved by standard numerical techniques. The jump conditions (\\ref{jump1}--\\ref{jump3}) provide  post-shock fluid variables. Given the  post-shock fluid variables, the equations (\\ref{a1}--\\ref{a4}) can be integrated inward to obtain the fluid variables inside the shock.  So far it seems that one can introduce the shock at any arbitrary radius. However, inward integration of the equations from an assumed shock position always end at a singular point, $\\lambda_0$, where at least one of the fluid variables becomes infinite. Since there is no natural way of eliminating this singularity, this must be the inner boundary of the baryonic shell, $M_{gas}(\\lambda_0)=0$. If the total energy of the system is to be conserved, the pressure at the inner boundary of the system must be zero, $P(\\lambda_0)=0$, otherwise, energy will continuously be injected into the system. It turns out that there is at most one value of $\\lambda_{\\rm s}$, for which the solution satisfies $P(\\lambda_0)=0$.  Smaller values of $\\lambda_{\\rm s}$ give positive mass and pressure at $\\lambda_0$, while larger values give zero mass and positive pressure. ",
        "conclusions": "For the gaussian initial density field with scale free power spectrum the rms density perturbation is $<(\\frac{\\delta M}{M})^2>\\propto r^{-(n+3)}$, where $-3<n<4$. It is therefore reasonable to identify  $s=(n+3)/2$, from which follows $0<s<3.5$. Thus in case $n>1$ we might expect that the universe contains a large number of spherical shells. We have shown in the paper that the evolution of spherical perturbations with $s>2$ results in formation of overdense regions, which, unless $(s-2)\\ll 1$ may contain a large fraction of matter.  A distinctive property of structure formation in the shell is the segregation between the baryons and the collisionless matter. The energy analysis shows that for $s<3.5$ (i.e. $n<4$), which follows from the gaussian density field, the collisionless matter cannot form bound objects, while at high redshifts cooling allows formation of the baryonic objects on the appropriate scales. Similarly, the explosive model (\\S 6) predicts that the baryonic shell is driven by pressure ahead of the collisionless and fragments in an independent way. This is entirely different from the evolution of positive perturbations, where gas must fall into the potential well formed by collisionless matter.   The energy analysis (\\S 5.1.2) shows that the mass of the bound objects formed from the baryonic shell is of order of Jeans mass at $T\\sim 10^4 K$ - $M\\sim 10^{11}\\Omega_{\\rm b}^{-1/2}(1+z)^{-3/2}M_\\odot$. This is again different from the evolution of positive perturbation, which predicts the Jeans mass to be also dependent on the spectrum of perturbations.  An observational limit on the scale and number of the shells can be obtained from the amplitude of Compton y-distortion. Both for the shell driven by the negative density perturbation and for the shell driven by internal pressure the predicted amplitude is of order $10^{-5}\\frac{\\Omega_{\\rm b}}{0.1}(\\frac{r_{10}}{3})^2f_F$, where $f_F$ is the present filling factor  (see Appendix B for the details of the calculation).  The current upper limit  $y=1.5\\cdot 10^{-5}$ is thus consistent with shells, whose present radius is below $\\sim 30$ \\rm {Mpc} and the filling factor is of order of unity."
    },
    "0201/astro-ph0201336_arXiv.txt": {
        "abstract": "We explore the prospects for using future supernova observations to probe the dark energy.  We focus on quintessence, an evolving scalar field that has been suggested as a candidate for the dark energy. After simulating the observations that would be expected from the proposed SuperNova / Acceleration Probe satellite (SNAP), we investigate two methods for extracting information about quintessence from such data.  First, by expanding the quintessence equation of state as $w_Q(z) = w_Q(0)-\\alpha\\ln(1+z)$, to fit the data, it is possible to reconstruct the quintessence potential for a wide range of smoothly varying potentials.  Second, it will be possible, to test the basic properties of the dark energy by constraining the parameters $\\Omega_Q$, $w_Q$ and $\\alpha$. We show that it may be possible, for example, to distinguish between quintessence and the cosmological constant in this way.  Further, when supernova data are combined with other planned cosmological observations, the precision of reconstructions and parameter constraints is significantly improved, allowing a wider range of dark energy models to be distinguished. ",
        "introduction": "There is now strong evidence, from observations of type Ia supernovae \\citep{SCP, Hi-z} and CMB anisotropy \\citep{Bernardis01,boom,dasi}, that we live in a universe that is geometrically flat and dominated by a nearly homogeneous component with negative pressure---the dark energy---which is causing the cosmic expansion to accelerate.  The existence of dark energy has recently been confirmed, independently of the supernova data, by combining the latest galaxy clustering data with CMB measurements \\citep*{GPEetal02, WTZ01}.  The most obvious dark energy candidate is vacuum energy, represented by the cosmological constant $\\Lambda$, which has pressure $p_\\Lambda = -\\rho_\\Lambda$.  Other dark energy candidates, include a network of topological defects and the much-discussed possibility of quintessence \\citep*{CDS98}, a spatially inhomogeneous, evolving component which is usually represented by a scalar field evolving in a potential (an idea first introduced by \\citet{PR88}).  In this paper we will focus on quintessence, though some of the results apply to other forms of dark energy  in so far as these can be parameterised by a simple equation of state.    It is our goal to explore what could be learned about the dark energy from future observations of type Ia supernovae (SNIa), such as may be possible with the proposed SuperNova / Acceleration Probe (SNAP) satellite \\citep{SNAP}.  This satellite aims to observe roughly 2000 supernovae a year for three years, with very precise magnitude measurements and negligible systematic errors, out to a redshift of $z=1.7$. Observations of this sort will permit a very precise measurement of the magnitude-redshift relation $m(z)$ and hence of the distance-redshift relation $r(z)$, which will probe the expansion history of the universe.   Current supernova observations already put weak constraints on the dark energy density $\\Omega_Q$ and equation of state $w_Q$. If the universe is assumed to be flat (\\emph{i.e.}, $\\Omega_k \\equiv 1-\\Omega_M-\\Omega_Q = 0$) and $w_Q$ is assumed to be constant in time, we have $\\Omega_Q\\ga 0.5$ and $w_Q\\la -0.4$ \\citep{SCP}.  These constraints improve considerably when they are combined with independent constraints from the CMB \\citep{GPE99} or large-scale structure \\citep*{PTW99}---most importantly, such combined data sets give $w_Q\\la -0.6$.   Several studies have examined the improved parameter constraints that may be possible with SNAP.  For example, \\cite{WA01,WA01-2} have found that SNAP can constrain a constant $w_Q$ to better than $10\\%$ accuracy, allowing some models to be distinguished from a cosmological constant ($w_\\Lambda=-1$).  In addition, many authors have examined the possibility of using SNAP to distinguish the evolution of $w_Q$ with redshift \\citep*{Astier00,BM00,Goliath01,HT00,MBS01,WA01,WA01-2}. There is a general consensus that SNAP data alone will not be able to distinguish an evolving equation of state.  Our results are in broad agreement with these studies.  It may also be possible to perform a direct reconstruction of $w_Q(z)$ and of the quintessence potential $V$ from supernova observations \\citep{HT99, CN00}. \\cite{Saini00} have attempted to perform reconstruction from the current supernova data and find, unsurprisingly, that nothing definitive can be learned at present. The theoretical studies that have been done have often encountered difficulty in accurately reconstructing the properties of a given model from simulated supernova data \\citep{HT00, WA01-2}.  This casts doubt upon the reliability of reconstruction.  One of our main goals in the present study is to develop a reliable method for producing accurate reconstructions from supernova data that is applicable to a wide class of quintessence models. Having done this, we will also want to explore whether such reconstructions will be useful.  The method we use for reconstruction demonstrates the power of using supernova observations to constrain cosmological parameters, and it highlights the usefulness of combining these observations with prior knowledge from other cosmological measurements. We therefore also undertake an exploration of the parameter constraints that will be possible with SNAP and its combination with other experiments, expanding on earlier analyses and discussing carefully what can and cannot be learned in this way.  This paper is arranged as follows. in Section~\\ref{sec:quint} we summarize the cosmological effects of quintessence pertaining  to supernova observations. In Section~\\ref{sec:arbfit} we simulate SNAP-like data for a quintessence-dominated universe and discuss the problems inherent in producing reconstructions from such data. We discuss a cosmologically parameterized fitting function for $r(z)$ in Section~\\ref{sec:cosfit}, and we use it for reconstruction in Section~\\ref{sec:cosrec}.    In Section~\\ref{sec:parest} we show how more general questions about the nature of the dark energy may be addressed by constraining parameters with SNIa data, and we discuss the long-term prospects for this sort of study in Section~\\ref{sec:ltparam}.  We draw conclusions about the prospects for observational tests of dark energy in Section~\\ref{sec:conclusion} ",
        "conclusions": "\\label{sec:conclusion}  In this study we have investigated the prospects for probing the dark energy, particularly quintessence, with observations of type Ia supernovae. We have examined tests for individual quintessence models (reconstruction) as well as tests that address more general questions about the dark energy (parameter estimation by likelihood analysis).  In both cases, we find that data from the proposed SNAP satellite will provide important information about dark energy, either alone or when combined with other cosmological observations.  In the case of reconstruction it will be important to take extreme care in choosing a function to fit SNAP's measurement of the distance-redshift relation $r(z)$. By using a fitting function based on a physical approximation to the coordinate distance, it is possible to produce accurate and reliable reconstructions from SNAP data combined with independent cosmological observations. Such reconstructions might provide useful observational constraints as we construct models for the dark energy.  To answer our basic questions about the nature of the dark energy, however, cosmological parameter estimation may be more useful. If the dark energy equation of state is significantly different from $-1$, the SNAP data alone will be able to rule out the cosmological constant.  Moreover, combining the SNAP constraints with independent measurements of $\\Omega_M$ will allow us to rule out $\\Lambda$ models for all but the most extreme cases of quintessence, when $w_Q$ is very close to $-1$. SNIa are less likely to provide evidence that $w_Q$ is evolving. Even combining SNAP data with very precise measurements of $\\Omega_M$ will only provide evidence of an evolving $w_Q$ in cases of extremely strong evolution.   Three-parameter fits may still be useful, however:  because they allow a more accurate determination of $w_Q(0)$ than two-parameter fits, they may be able to rule out a cosmological constant in cases where a two-parameter fit cannot do so.  The prospects are  good for probing dark energy with SNIa data. SNAP, by itself, may be able to constrain cosmological parameters with enough precision to rule out the cosmological constant as a dark energy candidate. However, it is when SNIa data are combined with other observations that their true value becomes apparent.  When combined with expected future measurements of $\\Omega_M$, for example, SNAP data provide much stronger constraints on the quintessence equation of state and may be able to produce useful reconstructions of the equation of state and quintessence potential. In each of these cases, the combination of several observations tells us much more than any of the observations by itself, and in each case, the SNIa observations are an essential component.  The measurements that may answer our questions about dark energy over the next decade or two constitute a difficult and impressive observational programme. Improved supernova observations are crucial to its success."
    },
    "0201/astro-ph0201100_arXiv.txt": {
        "abstract": "The effects of relativistic expansion on the late-time supernova light curves are investigated analytically, and a correction term to the (quasi-)exponential decay is obtained by expanding the observed flux in terms of \\(\\beta\\), where \\(\\beta \\) is the maximum velocity of the ejecta divided by the speed of light \\(c\\).  It is shown that the Doppler effect brightens the light curve owing to the delayed decay of radioactive nuclei as well as to the Lorentz boosting of the photon energies. The leading correction term is quadratic in \\(\\beta\\), thus being proportional to \\(E_{\\rm k}/(M_{\\rm ej} c^2)\\), where \\(E_{\\rm k}\\) and \\(M_{\\rm ej}\\) are the kinetic energy of explosion and the ejecta mass. It is also shown that the correction term evolves as a quadratic function of time since the explosion.  The relativistic effect is negligibly small at early phases, but becomes of considerable size at late phases.  In particular, for supernove having a very large energy(hypernova) or exploding in a jet-like or whatever non-spherical geometry, \\(^{56}\\)Ni is likely to be boosted to higher velocities and then we might see an appreciable change in flux. However, the actual size of deviation from the (quasi-)exponential decay will be uncertain, depending on other possible effects such as ionization freeze-out and contributions from other energy sources that power the light curve. ",
        "introduction": "It has long been recognized that the radiative transfer in supernova(SN) ejecta should be treated relativistically to account for the high velocities achieved in its outermost layers.  However, it is only in the last decade that the first numerical codes for radiative transfer that take relativistic effects into account were developed and applied to transfer problems in SN ejecta. At early times of explosion, the energy of photons emitted from the photosphere of SN ejecta is greatly enhanced by the Doppler effect.  As time goes on and the photosphere recedes towards the center of ejecta, the degree of the enhancement decreases, and eventually the resultant light curve(LC) is expected to follow an (quasi-)exponential deposition curve due to radioactive decays. However, there are several factors that cause the late-time LCs deviate from the radioactive decay curve.  One possibility is the flattening of LCs owing to the so-called ionization freeze-out effect, as first pointed out by \\citet{fra93} for the late-time LC of SN 1987A.  Another one is the signatures of contributions from other energy sources such as possible pulsar activity and circumstellar interaction as have been discussed by previous works.  In this Letter, I would like to point out that late-time SN LCs may show a deviation from radioactive decay curves as a result of the relativistic Doppler effect.  In particular, for very energetic SNe or what are called hypernovae(HNe), the relativistic effect would be very important since the maximum velocity near the surface of ejecta reaches a significant fraction of the speed of light.  It will be shown that the Doppler effect makes a LC brighter due to the Lorentz boosting, but at the same time the light-traveling-time effect and a pure-relativistic effect delay the decay of radioactive nuclei, thus tending to make the LC even brighter at late phases. The net result is determined by the sum of these two effects.  Late-time LCs provide direct information to determine the \\(^{56}\\)Ni masses ejected by SNe. Therefore, its precise determination is of great importance for studies of SN explosion mechanism and chemical evolution of galaxies.  In this Letter, we study the relativistic effect on the late-time LCs and estimate the size of the correction to the (quasi-)exponential decay for ordinary SNe and HNe.  For simplicity, we use the approximation that energy input from radioactivity is emitted as optical photons on the spot and carried away from the SN ejecta free of absorption. ",
        "conclusions": ""
    },
    "0201/astro-ph0201471_arXiv.txt": {
        "abstract": "{We present a calibration of the massive star formation rate vs. [CII] luminosity relation based on a sample of nearby, late-type galaxies observed with ISO-LWS and imaged in the H$\\alpha$+[NII] line. The relation holds for far-IR luminosities $10^8 \\leq L_{FIR} \\le 10^{10.5} L _\\odot$. The derived star formation rates have an uncertainty of about a factor of 10. Part of this uncertainty is due to the different mix of contributions to the [CII] emission from the different components of the interstellar medium in individual galaxies, as discussed in an appendix. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0201/astro-ph0201192_arXiv.txt": {
        "abstract": "{The relativistic iron line profile recently observed by {\\it XMM-Newton} in the spectrum of the Seyfert 1 galaxy MCG--6-30-15 (Wilms et al., 2001) is discussed in the framework of the {\\it lamp-post\\/} model. It is shown that the steep disc emissivity, the large line equivalent width and the amount of Compton reflection can be self-consistently reproduced in this scenario.}  \\authorrunning{A. Martocchia, G. Matt \\& V. Karas} \\titlerunning{On the broad, relativistic iron line of MCG--6-30-15 } ",
        "introduction": "Wilms et al. (2001; cited as W01 hereafter) recently presented and discussed an extremely broad and red-shifted iron K$\\alpha$ feature detected in the 06/11-12/2000 100 ksec {\\it XMM-Newton} observation of MCG--6-30-15. This Seyfert 1 galaxy is well known for possessing one of the best studied broad iron lines, whose profile is explained by relativistic effects (Tanaka \\ea, 1995, Guainazzi \\ea, 1998; see Fabian \\ea, 2000, for a review). The Fe \\ka\\ profile observed by {\\it XMM-Newton}'s {\\it EPIC-pn} camera is similar to the one observed by Iwasawa et al. (1996) using {\\it ASCA} data during a short ($\\sim 15.2$ ksec) period of low X-ray flux. W01 may have caught the source in a similar ``deep minimum state\", i.e. a state in which the primary flux is lower ($F_{\\rm 2-10 ~keV}=2.3 \\times 10^{-11}$ erg s$^{-1}$ cm$^{-2}$) and the line Equivalent Width (EW) higher (up to $300\\div400$ eV) than the time-averaged values.  The line profile indicates that a large fraction of the emission comes from $r<6r_{\\rm{g}}$ ($r_{\\rm{g}}=m={GM \\over c^2}$). This implies either that the central Black Hole (BH) is rotating -- thus the radius of the disc innermost stable orbit, $r_{\\rm{ms}}$, lies between $6r_{\\rm{g}}$ (=$r_{\\rm{ms}}$ for a static BH) and $1.23r_{\\rm{g}}$ ($r_{\\rm{ms}}$ of a canonically spinning BH: Thorne, 1974) -- or that the fluorescent line emission originates from matter falling freely below $r_{\\rm{ms}}$ (Reynolds \\& Begelman, 1997). Sako et al. (2001) proposed that also some spectral features observed at lower energies in this source, as well as in Mkn 766, are Ly$\\alpha$ lines of carbon, nitrogen, and oxygen, affected by relativistic broadening in the spacetime of a rotating BH.  In most works on the subject, a simple power law parameterization of the disc emissivity $\\epsilon(r) \\propto r^{-\\beta}$ is usually adopted. This is done also in the analysis of W01: letting $\\beta$ be a free fitting parameter, they find $\\beta \\sim 4$, a value much larger than usually found in Seyfert galaxies (Nandra et al., 1997). \\\\  In order to provide a physical picture of a so steep emissivity, W01 invoke strong magnetic stresses acting in the innermost part of the system, which dissipate a considerable amount of energy in the disc at very small radii. If the magnetic field lines thread the BH horizon, this would imply magnetic extraction of the BH rotational energy -- the so called {\\it Blandford-Znajek} effect (Blandford \\& Znajek, 1977; cited as BZ hereafter). However, the efficiency of the BZ effect has been questioned in recent years by e.g. Ghosh \\& Abramowicz (1997) and Livio, Ogilvie \\& Pringle (1999). These works argue that the electromagnetic output from the inner disc regions should in general dominate over that due to the BH. Thus the BH spin would probably be irrelevant to the expected electromagnetic power output from the system.  Krolik (1999), Agol \\& Krolik (2000) and Li (2000) proposed that MHD processes play a dominant role if magnetic field lines connect downfalling plasma near the hole with more distant regions: high efficiency of energy extraction can be achieved in this way even if the magnetic field does not thread the horizon itself. This magnetized accretion offers an alternative to the original BZ process and to its follow-up generalizations (e.g., Phinney 1983); however, the mechanism is violently non-stationary and such situations have not been quantitatively modelled yet (cf. Koide et al., 2000, and Tomimatsu \\& Takahashi, 2001, for the first attempts of such modelling). \\\\  In this paper we show that the required steep emissivity law, as well as the line EW and the amount of Compton reflection, may be reproduced with a phenomenological model in which a X-ray illuminating source is located on the BH symmetry axis ({\\it lamp-post} model: Martocchia \\& Matt, 1996, Petrucci \\& Henri, 1997, Bao, Wiita \\& Hadrava, 1998, Reynolds et al., 1999, Dabrowski \\& Lasenby, 2001). This can be considered as a simplified scheme, appropriate for various physical scenarios, including the mentioned MHD energy extraction. Indeed, Agol \\& Krolik (2000) state that magnetized accretion may also lead to enhanced coronal activity immediately above the plunging region. ``If so, this would provide a physical realization for models (...) which call for a source of hard X-rays on the system axis a few gravitational radii above the disc plane''. Alternatively, shock waves in an aborted jet close to the BH axis have been proposed as a source of the central irradiation by Henri \\& Petrucci (1997). This model assumes that a point source of relativistic leptons (e$^+$,e$^-$) illuminates the accretion disk by Inverse Compton process; the resulting angular and spectral distribution of soft and hard radiation has been derived. ",
        "conclusions": ""
    },
    "0201/astro-ph0201537_arXiv.txt": {
        "abstract": "{An analysis of H$\\alpha$ and H$\\beta$ spectra in a sample of 30 cool dwarf and subgiant stars is presented using MARCS model atmospheres based on the most recent calculations of the line opacities.  A detailed quantitative comparison of the solar flux spectra with model spectra shows that Balmer line profile shapes, and therefore the temperature structure in the line formation region, are best represented under the mixing length theory by any combination of a low mixing-length parameter $\\alpha$ and a low convective structure parameter $y$.  A slightly lower effective temperature is obtained for the sun than the accepted value, which we attribute to errors in models and line opacities. The programme stars span temperatures from 4800 to 7100~K and include a small number of population II stars.  Effective temperatures have been derived using a quantitative fitting method with a detailed error analysis.  Our temperatures find good agreement with those from the Infrared Flux Method (IRFM) near solar metallicity but show differences at low metallicity where the two available IRFM determinations themselves are in disagreement. Comparison with recent temperature determinations using Balmer lines by Fuhrmann~(\\cite{fuhrmann98, fuhrmann00}), who employed a different description of the wing absorption due to self-broadening, does not show the large differences predicted by Barklem~et~al.~(\\cite{bpo:hyd}).  In fact, perhaps fortuitously, reasonable agreement is found near solar metallicity, while we find significantly cooler temperatures for low metallicity stars of around solar temperature. ",
        "introduction": "Hydrogen is by far the most abundant species in typical stellar atmospheres.  In late-type atmospheres, a small fraction of hydrogen atoms capture free electrons forming H$^{-}$ ions, which dominate the continuum opacity.  Hydrogen itself is the main opacity source for earlier spectral types. The few lines that its simple atomic structure makes in the spectrum have a very distinct sensitivity to the atmospheric properties compared to metal lines. In optical stellar spectra, absorption lines of the Balmer ($n$=2) series are commonly used to study photospheres.  The well populated lowest levels of the atom produce considerable opacity at the centre of the lines, and interactions with charged ions, electrons and other hydrogen atoms result in extended wings in high-density atmospheres.  In late-type dwarfs, these wings are believed to form very close to LTE, in the deepest photospheric layers. As most protons are bound to electrons forming hydrogen, and hydrogen influences the main continuum opacity source, changes in the hydrogen abundance are hardly reflected in the lines' strengths, and the strengths are much more weakly affected by gravity and metal abundances than perturbations to the temperature. These properties attracted the attention of stellar spectroscopists, making Balmer lines a key feature in stellar classification schemes (e.g. Morgan et~al.~\\cite{morgan}).  Detailed analyses of Balmer lines in late-type stars are more recent (e.g. Gehren~\\cite{gehren1}; Fuhrmann et~al.~\\cite{fuhrmann93,fuhrmann94}; van't Veer-Menneret et~al.~\\cite{vbk}; Gardiner et~al.~\\cite{gardiner}). These modern studies exploit progress in theory and experiment on line broadening to infer stellar effective  temperatures from Balmer lines.   Vidal~et~al.~(\\cite{vcs:theory,vcs:tables}) developed a successful {\\it unified theory} to model the interaction of hydrogen atoms with charged particles.  Those calculations have been recently superseded by Stehl\\'e (\\cite{stehle94}) and Stehl\\'e \\& Hutcheon~(\\cite{stehle99}), who have computed Stark broadened line profiles including ion dynamic effects under the {\\it model microfield method}.  A further important broadening contributor is the collisions with neighbouring hydrogen atoms. Ali \\& Griem~(\\cite{ali_griem:errata}) used the multipole expansion of the resonance interaction potential in the impact approximation to calculate line-widths from this process. Barklem et al. (\\cite{bpo:let,bpo:hyd}, hereafter paper I and II respectively) have presented a self-broadening theory accounting also for dispersive-inductive interactions and without use of the multipole expansion. It was shown that the new description of the self-broadening would have a large impact on the computed Balmer line profiles, in particular when applied to derive effective temperatures for metal-poor dwarf stars.  Assuming a proper understanding of the line broadening of the hydrogen lines, the most notable difficulty for the use of their wings as a temperature indicator is the fact that they are formed in very deep layers. The thermal structure of the deepest optically transparent layers in late-type stars is significantly affected by convection.  Simple modelling of surface convection is still a challenge. The commonly used mixing-length formalism incorporates unphysical parameters which are hard to connect with quantities that can be derived from observations or hydrodynamical simulations and, therefore, are difficult to constrain. This obstacle makes flux-constant homogeneous models particularly uncertain in these layers. Other theories that dispense with the free parameters in the mixing-length theory (MLT) have also been proposed (e.g. Canuto et al. \\cite{canuto1}; Canuto \\& Mazzitelli \\cite{canuto2}); however, Gardiner et al. (\\cite{gardiner}) found that after adjusting the mixing-length parameter $\\alpha$, MLT  performed similarly. Semi-empirical modelling (e.g. Allende Prieto et al.~\\cite{inversion}) does not offer a viable solution, as the employed metal lines do not probe layers as deep as those where hydrogen lines are formed. These theoretical problems related to establishing a model atmosphere combine with observational constraints. Balmer lines require high dispersion observations with  a large spectral coverage, and a predictable instrumental response, making possible a methodical and accurate continuum normalisation. In this situation, use of Balmer lines as a part of spectroscopic analyses requires a careful assessment of all possible sources of error.  In this paper we investigate the impact of the aforementioned line broadening theory advances in the framework of 1D model atmospheres, and attempt to identify those areas where improvement is most desirable.  In Sect.~2 we describe our observations and reduction procedure.  In Sect.~3 we describe how model spectra are computed, and put the atmospheric models in context with others.  In Sect.~4 we introduce a method for quantitative comparison of observations and model spectra, and subsequently an automated fitting procedure for deriving effective temperatures. In Sect.~5 we make a detailed survey of possible errors and their effect on effective temperatures.  Application to the solar spectrum and the programme stars is then presented.  Finally in Sect.~6 the results are compared with other work, and in Sect.~7 our conclusions are presented. ",
        "conclusions": "Effective temperatures have been derived from Balmer line profiles using a quantitative fitting method with a detailed error analysis including investigation of the susceptibility to various errors with temperature and metallicity for our models.  Our temperatures find good agreement with the IRFM near solar metallicity but show differences at low metallicity where the two available IRFM determinations themselves are in disagreement.  Our results for metal-poor stars seem in better agreement with Magain~(\\cite{magain}), though we caution this is only for 4 stars and the better agreement may be fortuitous.  This should be investigated further, for which more, and better observations of Balmer lines in population II stars are needed.   The origin of the differences between these two IRFM determinations needs also to be understood.  Through estimates of the uncertainties, we found that the relative weight that should be given to H$\\alpha$ and H$\\beta$ \\emph{in determining effective temperatures} varies quite substantially with stellar parameters.  This is predominantly a matter of balance between the reduced sensitivity of H$\\alpha$ to temperature at low metallicities making broadening, $\\log g$ and observational uncertainties very important, and the high sensitivity of H$\\beta$ to convection.  Thus to improve the accuracy of Balmer line temperatures for metal-poor stars these four areas must be addressed.  Observational uncertainties will continue to improve as large telescopes provide the possibility of combined high resolution and SNR for these faint objects.  There no doubt exists a limiting accuracy for continuum determination procedures for echelle spectra such as that used here, but we believe this has not yet been reached.  Fuhrmann~et~al.~(\\cite{fuhrmann94}) suggest an accuracy of about 0.3$\\%$ is achievable with modern spectroscopy.  Planned satellite astrometry missions should address the gravities.  Both the remaining uncertainties, convection and self-broadening, lie in the theoretical realm at least at present.  Considerable progress has been made in 3D hydrodynamical simulations of convection and inhomogeneities in cool stars, and this will be the way forward.  However, for the immediate future where direct application of such models is impractical, calibrations of MLT parameters \\emph{for Balmer lines} across the HR diagram similar to Ludwig~et~al.~(\\cite{ludwig}) or tabulated $T_\\mathrm{eff}$ corrections (against a given MLT parameter set) would be important.  Improved calculations of the self-broadening of Balmer lines should be undertaken without resort to the impact approximation which will require improved short range potentials from those used in paper II.  In paper I and II large differences between $T_\\mathrm{eff}$ values derived with Ali \\& Griem~(\\cite{ali_griem:errata}) theory and new calculations of the self-broadening, particularly in metal-poor stars, were predicted.  Comparison of our results with those of Fuhrmann~(\\cite{fuhrmann98,fuhrmann00}) where Ali \\& Griem theory has been employed, find in fact good agreement except for low metallicity stars of around solar temperature where the differences are smaller than expected, but the temperatures are still significantly cooler. The reason for this unexpected agreement is a number of systematic differences in models and employed broadening recipes, which though typically individually small, together somewhat compensate the difference made by the new self-broadening. This emphasises that in order to achieve high absolute precision in $T_\\mathrm{eff}$ determinations from Balmer lines such small differences must in fact be carefully considered.  Errors in the relative temperature scale, particularly between stars of low and solar metallicity important in tracing chemical evolution in the galaxy, are unlikely to be much smaller than the errors in the absolute scale since those errors which might cancel in differential comparison, namely broadening theory and model errors (including convection), typically vary with metallicity resulting in little cancellation."
    },
    "0201/astro-ph0201067_arXiv.txt": {
        "abstract": "{ We discuss the morphology and spectrophotometry of 5 comets visible in August, 2001. We decompose comae into coma profiles and azimuthally renormalized images, in which general and local features are quantitatively comparable. Comet 19P/Borrelly showed a strong gas fan toward the solar direction, but no detectable gas in the tail. Dust in its inner coma was collimated toward the antisolar direction and the tail, with no dust in the outer coma. The contribution of spatial variations structure was moderate, about 35\\%. Comet 29P Schwassmann-Wachmann 1 was observed in outburst: we detected ``spinning'' jet structures. A high level of dust production resulted in an unusually high $Af{\\rho}$=16600 cm. The spatial variations reached $-$77\\%, at the minimum, due in part to a jet and a ring-like structure in 1 arcminute distance from the nucleus. In comet C/2001 A2, we detected a strong post-perihelion increase of dust and gas activity, in which the C$_2$ profile became one magnitude brighter over a 3-day period. For comets C/2000 SV74 and C/2000 WM1, we present detailed pre-perihelion spectrophotometry and morphological information. Comet C/2000 SV74 showed high dust production ($Af{\\rho}=1479$ cm). Its coma suggests a steady-state outflow of material, while the low contribution of spatial variations support high level activity. The coma of C/2000 WM1 is dominated by solar effects, and CO+ forms the bulk of its gas activity. Despite its large heliocentric distance, we observed a nice tail. ",
        "introduction": "The distribution of distinctive cometary components (dust, ions, radicals) can be well studied with help of two-dimensional, narrow-band CCD images. In addition to measuring the spatial distribution of species, one can measure column densities of gas and dust, and draw conclusions on chemical composition, production rates and activity levels. Despite the efficiency of this method, only a small fraction of visible comets are studied with detailed spectrophotometry. Spectroscopic methods, on the other hand, have the advantage of producing a spectrum with a well-defined continuum level, so emission features show clearly. Large spectrophotometric and spectroscopic surveys of many comets were published by A'Hearn et al. (1995) (hereafter A95) and Fink \\& Hicks (1996), where the interested reader can find detailed descriptions of methods.  Surface photometry of images can address coma morphology (such as radial coma profiles and non-radial features, also called as coma profiles and azimuthally renormalized images, see, e.g., Lederer et al. 1997, Larson \\& Slaughter 1991), and also may yield estimated nuclear radii (e.g. Luu \\& Jewitt 1992, Lamy \\& T\\'oth 1995, Lowry et al. 1999). An appropriate selection of medium- and narrow-band filters centered on different wavelengths can separate the dust continuum from emission by gas. Differences in gas and dust components reveal the effect of radiation pressure on different types of particles (a good example of combined quantitative coma analysis can be found, e.g. in Schulz et al. 1993).  The main aim of our work is to contribute to this field of ground-based solar system research with new narrow- and medium-band spectrophotometric observations of comets visible in August, 2001. The function of this paper is to present the results of observations of 5 comets carried out at Calar Alto Observatory. The paper is organised as follows. Sect.\\ 2 deals with the observations, methods of analysis are described in Sect.\\ 3, results on individual objects are given in Sect.\\ 4, while the discussion is presented in Sect.\\ 5. This work is an extension of our previous analysis of distant active comets (Szab\\'o et al. 2001). ",
        "conclusions": "Five comets were observed few days before new moon in August, 2001. The comet producing the highest $Af{\\rho}$ was 29P/Schwassmann--Wachmann 1 as it was caught during its outburst. Other comets also showed unusually high $Af{\\rho}$ values, especially C/2000 SV74, which is predicted to become a fairly bright, dust-rich comet at perihelion.  Inner comae of comets often show more matter ejected to the solar side than elsewhere. This asymmetry of matter production is probably supported by the warm solar side of the nucleus. Comparing the continuum-dominated images and the emission-dominated ones, we find a significantly changing coma composition for C/2001 A2 in a smaller outburst. In the case of 19P and C/2001 A2, the solar side is dominated by gaseous components, due to interaction with the radiation pressure and solar wind: dusty components are blown backwards and the solar side becomes gas-rich. Outbursts, jets or proximity to the Sun may enrich gas in comae, and may lead to anomalous coma composition. We believe these factors explain the varying emission surface brightness profiles of C/2001 A2 in outburst.  \\begin{table} \\caption{Maxima and minima of spatial variations characterized as $s$ solar and $a$ antisolar, $r$ radial or $t$ tangential-type peaks. Distance from the nucleus is shown in $10^3$ km.} \\label{peak} \\begin{center} \\begin{tabular} {lllrlrl} \\hline Ref. & R(AU) & $\\Delta$(AU) & $D_{\\rm max}$ & $lc_{\\rm max}$ & $D_{\\rm min}$ & $lc_{\\rm min}$\\\\ \\hline 19P & 1.400 &1.685 & 5.5 & 0.342$s$ & 3.7 & $-$0.398$t$ \\\\ 29P & 5.919 & 5.121 & 26.1 & 0.370$s$ & 18.6 & $-$0.769$t$ \\\\ SV74 & 4.244 & 4.036 & 11.7 & 0.331$a$ & 5.9 & $-$0.218$rs$ \\\\ WM1 & 2.798 & 2.891 & 10.5 & 0.423$a$ & 7.3 & $-$0.603$rs$\\\\ A2a & 1.628 & 0.714 & 2.3 & 0.224$a$ & 1.3 & $-$0.217$rs$\\\\ A2c & 1.669 & 0.764 & 2.2 & 0.159$a$ & 1.4 & $-$0.151$rs$\\\\ \\end{tabular} \\end{center} \\end{table}  In order to compare quantitatively the asymmetries in comets, the contributions of spatial variations are compared in Table\\ 5. Peaks of maximal intensity are distinguished by their location relative to the center of the coma: $s$ means peaks on the solar side, $a$ denotes peaks on the anti-solar side. We classify peaks of minima as follows: $r$ denotes radial-type peaks appearing near the radial section (solar or antisolar type), $t$ denotes tangential-type peaks near the tangential section. The local contribution of spatial variations at minimum peaks seem to be well correlated with the geocentric radius. Similar but less obvious correlations can be found for the maxima, too. Generally, peaks of the azimuthally renormailzed images develop farther from the nucleus, and their contribution to the total coma intensity increases with increasing solar distance. This conclusion is consistent with the view that solar wind and radiation pressure affects the cold and slowly outflowing coma of less active comets at a larger solar distance.  C/2000 SV74 seems to be an exception to the rule, as its coma is significantly (by a factor of 3) less affected by spatial variations than the other four comets. This character can be hardly explained by the small phase angle or the generally circular appearance, and supports the idea that high level matter production is present that can challenge the solar wind.  Dependencies between Gunn-colors and  $Af{\\rho}$ parameter have been found for comets. Below we summarize the correlation, their standard errors (in brackets) and their regression coefficient. Independent correlations with better regression coefficient than 0.80 or $-$0.80 have been accepted.  \\begin{tabular} {ll} \\\\ $(v-g)=-0.51(20)+0.214(7)\\cdot {\\rm log}~Af{\\rho}[cm]$ & \\\\ \\hskip1cm regr. coeff.=0.87 &\\\\ $(r-j)=0.12(5)+0.7(3)\\cdot (v-g)$ & \\\\ \\hskip1cm regr. coeff.=0.81 & \\\\ $(j-z)=0.19(4)-0.8(2)\\cdot (r-j)$ & \\\\ \\hskip1cm regr. coeff.=-0.88 &\\\\ \\\\ \\end{tabular}  We explain the correlation between $v-g$ color and $Af{\\rho}$ as a scattering effect of dust. Matter particles scatter the short-wavelength violet color the most, thus, the more dust present in the coma, the more \"reddish\" $(v-g)$ its color. Continuum colors of 29P/Schwassmann-Wachmann 1 are measured to be slightly less reddish than previously published values: $B-V=0.8$ (Hartmann et al., 1982), $V-R=0.502$ and $R-I=0.492$ (Meech et al. 1993), which may be transformed into the Gunn-system with help of the transformations determined by Kent (1985) and J\\o rgensen (1996), yielding $v-g=0.478$, $g-r=0\\fm028$. Meech et al. (1993) attribute reddish colors to a large distance from the Sun. In present paper, the experimental correlation between $Af{\\rho}$ and (v-g) also explains quite reddish colors.  Correlation between $r-j$ and $j-z$ may be an artificial effect. Examining the individual comets, the $j-z$ color is found to correlate with solar distance as the radiation contribution in the near-infrared has increasing effect with decreasing solar distance.  In order to understand production rates, we have compared our data to the extensive set of A95. In the case of 19P/Borrelly, production rate ratios are in a perfect agreement with the previous results, as mentioned in A95. The detected lance-head shape of the coma is similar to the HST images taken by Lamy et al. (1998). Their slope parameter varying around $-$1 with moderate angular variations is in a perfect agreement with our observations. We note that Schleicher (2001b) published $Af{\\rho}$ measurements in September, which were half those measured by us.  The slope parameter $G=-1.9$ of Schleicher (2001b) is also not comparable with our smooth profile.  C/2000 SV74 and C/2000 WM1 have similar appearances, although during their further evolution significant structural differences may develop, due to the difference in their perihelion distances (3.54 and 0.56 AU). The abundance of C$_2$ in C/2000 SV74 is much higher than usual: one can find only 3 comets (P/Russel 4, C/Shoemaker 1984 XII, C/Shoemaker 1984 XV) with higher C$_2$/CN production among the 85 objects discussed in A95.  C/2001 A2 showed a significant change in composition during its outburst. Compared to the statistics and Fig.\\ 4c of A95, the [$Af{\\rho}$]-[CN] seems to be anomalously high in ``quiescent'' state ($-22.07$), while is also a bit high, but not peculiar, in ``outburst'' ($-22.74$); these values fall within the $-23.1\\pm0.9$ interval for all of the comets discussed in A95. As a possible explanation, one can imagine that C/2001 A2 was in a period of anomalously low activity, a ``negative outburst'', during the time of high [$Af{\\rho}$]-[CN] production (e.g. A2a, on 13th August, 0:14 UT). The normal behaviour in outburst-like events is to show increased gas activity.  Comets 19P and C/2001 A2 show similarly bluish colors and CN-rich gas production. Low [C$_2$]-[CN] rates allow us to classify C/2001 A2 as a Borrelly-type comet (Fink et al. 1999). Note that the drastic variations observed during some months previously might also influence the classification.  Finally, we summarize our results as follows.  \\noindent 1. The use of Gunn photometric system in structural and production analysis of comets is demonstrated. When augmented with comet and continuum interference filters, the system combines the advantages of the two systems usually preferred in cometary astronomy (Johnson filters for morphology and narrow-band filters for spectrophotometric studies).  \\noindent 2. A technique for generating and analyzing azimuthally renormalized images or non-radial residual maps is discussed, and the power of this tool to study morphology is demonstrated. In the case of C/2000 SV74, C/2000 WM1 and C/2001 A2, simple comae with slight solar formation effects are observed. In the case of 19P/Borrelly, slope parameter was almost $-1$, but spatial variations were observed to make moderate contributions. 29P/Schwassmann-Wachmann 1 was observed in outburst, and a well-developed jet and a ring at larger nuclear distance were detected. Variations from spherical outflow are indicated by local contribution parameters at the peaks of the azimuthally renormalized images. Generally, the contribution of the non-radial parts to the whole coma was found to increase with the solar distance.  \\noindent 3. C/2000 SV74 and C/2000 WM1 were observed a few months before their perihelion.  Large $Af{\\rho}$ values, the extended coma of C/2000 SV74 and the nice tail of C/2000 WM1 suggest that they will be interesting objects in perihelion. C/2000 SV74 showed unusually high [C$_2$]-[CN] production ratio, and developed a nearly spherical coma at a distance of 4 AU. During our observations, C/2001 A2 suffered a smaller outburst, ejecting mainly gaseous components.  The ejected matter evolved in a spherically symmetric manner. The amount of gas increased by a factor of 4 during the outburst, while the dust components increased by about 39\\%. Its matter production is quite atypical when low level of activity is present."
    },
    "0201/astro-ph0201251_arXiv.txt": {
        "abstract": "{ We present CCD UBVI observations obtained in the field of the two previously unstudied dissolving open cluster candidates NGC~7036 and NGC~7772. Our analysis suggests that  both the objects are Open Cluster Remnants (OCR).\\\\ NGC~7036 is an open cluster remnant with a core radius of about 3-4 arcmin. We derive for the first time estimates of its fundamental parameters. We identify 17 likely members which define a group  of stars at 1 kpc from the Sun, with a low reddening E$(B-V) \\approx 0.1$, and with an age of about 3-4 Gyr.\\\\ As for NGC~7772, we identify 14 likely members, which define a group of stars with a very low reddening (E$(B-V) \\approx 0.03$), 1.5 Gyr old and located about 1.5 kpc from the Sun. ",
        "introduction": "The dynamical evolution and the final fate of open star clusters in the Milky Way is nowadays a very active research field. Open star clusters are weakly bound objects with a typical lifetime of less than a Gyr (Dutra \\& Bica 2000, Bergond et al. 2001), which ultimately depends on the initial mass of the cluster, the birthplace and the fraction of primordial binaries (de la Fuente Marcos 1998, 2001).\\\\ Recently, Bica et al. (2001) draw the attention on a sample of high Galactic latitude ($b > 15^{o}.$) star clusters presumably in a advanced stage of dynamical evolution, which they baptized {\\it Probable Open Cluster Remants (POCR)}. The prototype of this class of objects is NGC~6994 (M~73), recently studied by Bassino  et al. (2001) and Carraro (2001), who performed the first multicolor photometric studies of this cluster, but arrived at opposite conclusions on the nature of this object. While Bassino et al. suggest that M~73 is the remnant of a star cluster, Carraro (2000) proposes that it is just a chance alignment of four bright stars.\\\\ Although it is a difficult task to unravel the nature of a star concentration basing only on photometry, it however remains the first necessary step. Indeed sometimes the Color Magnitude Diagrams (CMDs) and Color-Color Diagrams (CCDs) are sufficient to disentangle between a real bound system or a random enhancement of stars (Carraro \\& Patat 1995, Piatti \\& Clari\\`a 2001). In many cases however, photometry cannot help to decide unambiguously about the nature of a star concentration: in this situation radial velocities and/or proper motions studies are necessary (Baumgardt 1998, Baumgardt et al. 2000, Odenkirchen \\& Soubiran 2002).\\\\ The census provided by Bica  et al. (2001) lists 20 candidate dissolving open clusters, some of which completely unstudied. This is the case of NGC~7036 and NGC~7772, two high latitude objects traditionally  considered genuine open clusters, which are the subject of the present study.\\\\ The basic idea of this paper is to present the first photometric study of these clusters and to provide a list of probable members stars to be further studied with high resolution spectroscopy.\\\\ A nice example of that is the recent spectroscopic follow-up of NGC~6994 by Odenkirchen \\& Soubiran (2002), who confirmed Carraro (2000) suggestions that this object is a chance alignment of four bright stars.\\\\    In Sect.~2 we briefly present the observations and data reduction. Sects.~3 and 4 illustrate our results for NGC~7036 and NGC~7772, and, finally, Sect.~5 draws some conclusions and suggests further lines of research.     \\begin{table} \\caption{Basic parameters of the observed objects. Coordinates are for J2000.0 equinox} \\begin{tabular}{ccccc} \\hline \\hline \\multicolumn{1}{c}{Name} & \\multicolumn{1}{c}{$\\alpha$}  & \\multicolumn{1}{c}{$\\delta$}  & \\multicolumn{1}{c}{$l$} & \\multicolumn{1}{c}{$b$} \\\\ \\hline & $hh:mm:ss$ & $^{o}$~:~$^{\\prime}$~:~$^{\\prime\\prime}$ & $^{o}$ & $^{o}$ \\\\ \\hline NGC~7036       & 21:10:02 & +15:31:06 & 64.55  & -21.44\\\\ NGC~7772       & 23:51:46 & +16:14:48 & 102.74 & -44.27\\\\ \\hline\\hline \\end{tabular} \\end{table} ",
        "conclusions": "We have presented the first CCD $UBVI$ observation of the Probable Open Cluster Remnants NGC~7036 and NGC~7772.\\\\ Our analysis suggests that:   \\begin{itemize}  \\item NGC~7036 and NGC~7772 are two nice examples of intermediate-age open clusters in advanced stages of dynamical evolution; \\item NGC~7036 is a 3-4 Gyr open cluster located 1 kpc far from the Sun. However, we stress the fact that this cluster still remains a somewhat doubtful  case, due to the absence of clear features in the CMDs; \\item NGC~7772 is a less doubtful case. We showed that the cluster underwent strong low mass stars depletion. What remains is a group of 14 stars 1.5 Gyr old and located 1.5 kpc away from the Sun. \\end{itemize}  \\noindent It is worth to remarking that the present results must be considered with some caution, and that the list of members must be better constrained by determining individual star radial velocities and proper motions. This way these objects can become templates for N-body simulation aimed at investigating the dynamical evolution of open star clusters and the origin of the field star population"
    },
    "0201/astro-ph0201298_arXiv.txt": {
        "abstract": "We present an analysis of interstellar absorption along the line of sight to the nearby white dwarf star \\HZ. The distance to this star is 68$\\pm$13 pc, and the line of sight extends toward the north Galactic pole. Column densities of \\Oone, \\None, and \\Ntwo\\ were derived from spectra obtained by the Far Ultraviolet Spectroscopic Explorer (\\fuse), the column density of \\Done\\ was derived from a combination of our \\fuse\\ spectra and an archival \\hst\\ GHRS spectrum, and the column density of \\Hone\\ was derived from a combination of the GHRS spectrum and values derived from \\euve\\ data obtained from the literature. We find the following abundance ratios (with \\twosig\\ uncertainties): \\Done/\\Hone\\ = $(1.66 \\pm 0.28)\\times 10^{-5}$, \\Oone/\\Hone\\ = $(3.63 \\pm 0.84)\\times 10^{-4}$, and \\None/\\Hone\\ = $(3.80 \\pm 0.74)\\times 10^{-5}$. The \\Ntwo\\ column density was slightly greater than that of \\None, indicating that ionization corrections are important when deriving nitrogen abundances. Other interstellar species detected along the line of sight were \\Ctwo, \\Cthree, \\Osix, \\Sitwo, \\Arone, \\Mgtwo, and \\Fetwo; an upper limit was determined for \\Nthree.  No elements other than \\Hone\\ were detected in the stellar photosphere. ",
        "introduction": "Deuterium is one of the primary products of big bang nucleosynthesis (BBN), and its primordial abundance provides a sensitive measure of the baryon density of the universe \\citep{Schramm:1998, Burles:2000}. However, deuterium is consumed in stars far more quickly than it is produced by the P-P cycle or by any other known stellar nucleosynthetic process \\citep{Epstein:1976}; thus, measurements of the deuterium abundance at the present epoch can only provide lower limits to the primordial abundance.  Measurements obtained at a variety of redshifts will sample the chemical evolution of the universe, and may permit extrapolation to a value for the primordial abundance of deuterium. Measurements of the abundance of deuterium and various products of stellar nucleosynthesis, such as oxygen, within our Galaxy will improve our understanding of chemical evolution as well as provide a lower limit to the primordial deuterium abundance.  Recent reviews of deuterium abundance measurements can be found in \\citet{Lemoine:1999}, \\citet{Linsky:1998}, and \\citet{Moos:2001}.  \\citet{Linsky:1998} reviews measurements of the deuterium abundance for material in the local interstellar cloud (LIC), which are consistent with a common value for D/H of $(1.5\\,\\pm\\,0.1) \\times\\,10^{-5}$.  The LIC and other nearby clouds are thought to be embedded in a large bubble of hot, low density gas ($T\\approx\\,10^{6}K,\\ n\\approx\\,0.005\\,cm^{-3}$) that was probably produced by supernovae and stellar winds arising in the Scorpius-Centaurus OB association \\citep{Cox:1987, Frisch:1995}.  If this picture is correct, then the local interstellar medium (LISM) may contain an inhomogeneous mixture of clumps of older material swept up by the expanding bubble along with the more-recently processed material in the bubble itself.  Thus, even in this local environment one may find regions of gas with somewhat different evolutionary histories and different compositions. The same processes have been at work throughout the history of the Galaxy; hence measurements of deuterium and other abundances along numerous lines of sight will help determine not only the chemical evolution of the Galaxy, but also the relative timescales for chemical evolution and the mixing of material.  In this paper we present an analysis of the line of sight to the nearby white dwarf \\HZ.  Companion papers will present similar analyses of the lines of sight to G\\,191-B2B \\citep{Lemoine:2001}, WD\\,0621-376 \\citep{Lehner:2001}, WD\\,1634-573 \\citep{Wood:2001}, WD\\,2211-495 \\citep{Hebrard:2001}, BD\\,$+28^{\\degr}~4211$ \\citep{Sonneborn:2001}, and Feige\\,110 \\citep{Friedman:2001}. An overview will be provided by \\citet{Moos:2001}.  High-resolution spectra of the DA white dwarf \\HZ, covering the far ultraviolet (FUV) wavelength range 905--1187\\AA\\ were obtained with the {\\it Far Ultraviolet Spectroscopic Explorer} (\\fuse) for the purpose of studying the deuterium abundance of the local interstellar medium (ISM).  This line of sight is promising for a study of D/H for several reasons: the star itself is moderately bright, exhibits a nearly featureless continuum, and has been accurately modelled; the absorbing interstellar gas appears to have only a single velocity component and the column density is very low, so many absorption lines that are ordinarily on the flat part of the curve of growth are either unsaturated or exhibit only mild saturation effects.  In addition, the \\Hone\\ column density along the line of sight can be determined by two independent methods: from the shape of the Lyman continuum observed by the {\\it Extreme Ultraviolet Explorer} (\\euve), and from the \\LA\\ profile observed with the Goddard High Resolution Spectrograph (GHRS) onboard the Hubble Space Telescope (\\hst).  Systematic uncertainties in determination of \\Hone\\ column densities are often the limiting factor in measurements of the D/H ratio, so the availability of multiple independent measurements of N(H) is quite valuable.  In Section 2 we discuss the line of sight to \\HZ\\ and the properties of the star; in Section 3 we describe the observations and data reduction procedures; in Section 4 we describe our analysis procedures, our methods for minimizing systematic errors, and our measured column densities; and in Section 5 we summarize our results. ",
        "conclusions": "The line of sight to \\HZ\\ was investigated using a combination of \\fuse\\ spectra, archival GHRS spectra, and results obtained from \\euve\\ spectra taken from the literature.  The line of sight to \\HZ\\ is particularly simple, with only a single velocity component discernible.  An additional high-temperature component was found with a very low column density, detectable only in \\Hone.  The line of sight to \\HZ\\ exits the LIC after a very short distance, and the velocity of the main component of the absorbing gas is not compatible with the velocity of the nearby G cloud in the LISM, hence this material is presumed to be located in the North Galactic Pole cloud. Doppler widths for each species determined from profile fits or curve of growth analyses of the \\fuse\\ data were consistent with the temperature and turbulent velocity for the gas derived from the GHRS high resolution data.  Consistent fits to the \\Done\\ column density were obtained when fitting the \\fuse\\ and GHRS data separately and simultaneously, indicating that the \\Done\\ \\LA\\ line is not affected by saturation effects.  The \\fuse\\ spectra were used to determine \\Oone\\ and \\Done\\ column densities by both profile fitting and curve of growth analyses, with consistent results obtained in each case. Careful attention was paid to numerous potential sources of systematic error.  Our adopted \\Hone\\ column density, \\lognh = 17.93 $\\pm$ 0.06 (\\twosig), is the mean of our measurement from the GHRS \\LA\\ profile and the EUV measurements of \\citet{Dupuis:1995}, \\citet{Barstow:1997}, and \\citet{Wolff:1999}.  Our results are summarized in Table \\ref{tab_summary}. All uncertainties in the table are \\twosig\\ (but values quoted below from the literature are \\onesig). Our result for D/H along the line of sight to \\HZ\\ is $(1.66 \\pm 0.28)\\times 10^{-5}$. The GHRS dataset analyzed in this work has been previously analyzed by \\citet{Landsman:1996}, who obtained the similar value for D/H of $1.6 \\times 10^{-5}$. Our value for D/H along this sightline exceeds by approximately 1.5\\,$\\sigma$\\ both the mean value reported by \\citet{Linsky:1998} for measurements of D/H within 100\\,pc of the sun, $(1.47 \\pm\\ 0.10) \\times 10^{-5}$, and the mean value of the \\fuse\\ measurements reported in this first set of papers \\citep{Moos:2001}.   Our result for \\Oone/\\Hone\\ is $(3.63 \\pm\\,0.84)\\times 10^{-4}$. Converting this ratio to a logarithmic abundance log~(\\Oone/\\Hone)~+~12.00 gives 8.56\\err{0.09}{0.11}.  The ionization balance for both O and H are linked by resonant charge exchange reactions (see \\citealt{Jenkins:2000} and references therein), and there is no apparent \\Hmol\\ present, so \\Oone/\\Hone\\ should be representative of O/H along this line of sight. \\citet{Meyer:1998} report an average gas-phase abundance for 13 sight lines of $(3.43 \\pm\\,0.15)\\times 10^{-4}$ \\citep[after correction for our preferred \\Oone\\ $\\lambda1355$ oscillator strength of $f = 1.16\\times10^{-6}$ from ][]{Welty:1999}, and they estimate no more than O/H$\\sim1.8\\times 10^{-4}$ can be incorporated into grains, implying an upper limit to the total (gas+dust) ISM oxygen abundance of $5.2\\times 10^{-4}$.  Our O/H measurement is consistent with the average gas-phase value of \\citet{Meyer:1998}.  \\citet{Sofia:2001a,Sofia:2001b} have demonstrated that the value of O/H determined from F and G star photospheres, $(4.45 \\pm 1.56)\\times 10^{-4}$, is consistent with two recent determinations of the solar abundance, $(5.45 \\pm 1.00)\\times 10^{-4}$ \\citep{Holweger:2001} and $(4.90 \\err{0.60}{0.53})\\times 10^{-4}$ \\citep{AllendePrieto:2001}, and with the  total oxygen abundance estimated from the work of \\citet{Meyer:1998}. \\citet{Sofia:2001a,Sofia:2001b} argue that the recent solar system values should therefore be used as the ISM abundance standard for oxygen.  Our value for O/H in the gas phase along the line of sight to \\HZ\\ is significantly lower than each of these values, although not by more than \\twosig\\ (given the large uncertainties in the solar system determinations).  Our result is in agreement with the compilation of B-star abundances presented by \\citet{Sofia:2001a,Sofia:2001b}, $(3.50 \\pm 1.33)\\times 10^{-4}$, although these authors argue that B-star abundances are unlikely to be appropriate measures of the ISM standard.  It seems likely that the gas-phase oxygen abundance measured by \\citet{Meyer:1998} implies modest depletion of oxygen in the ISM towards more distant stars; given the agreement of our O/H measurement with the Meyer et al. average, it seems likely a similar degree of oxygen depletion is present in the cloud observed along the \\HZ\\ sight line.    The ionization fractions for N and H are not as strongly coupled as those of O and H \\citep{Jenkins:2000}, so it is possible for N to be more highly ionized than H, and therefore for the gas-phase abundance of N relative to its solar system value to be less than that of O.  If we combine our column densities for \\None, \\Ntwo, and the maximum likely column density for \\Nthree\\ (assuming only half of the measured absorption is due to \\Sitwo), we obtain a total nitrogen column density of $N_N = 7.88\\err{1.56}{1.43} \\times 10^{13}$, and a ratio N/\\Hone\\ = $(9.26 \\pm 1.87)\\times 10^{-5}$.  This ratio provides an upper limit to N/H, corresponding to the limit that that hydrogen along this line of sight is predominantly neutral.  Given the significant column density of \\ion{N}{2}, it is reasonable to expect that there will also be a significant column density of \\ion{H}{2} present.  Converting to a logarithmic scale with \\lognh\\ = 12.00 gives an abundance for N of 7.97\\err{0.08}{0.10}.  This is consistent with the solar abundance value of \\citet{Grevesse:1993}, 7.97, and the more recent value from \\citet{Holweger:2001} of 7.93 $\\pm$ 0.11, and somewhat exceeds the value of 7.81 for nearby B stars \\citep{Gies:1992}. We can set a lower limit to the N abundance by setting the ionization fraction of H to the upper limit of 50\\% found by \\citet{Dupuis:1995} (the other \\euve\\ analyses obtained lower ionization fractions for hydrogen).  The resulting logarithmic abundance is 7.79\\err{0.08}{0.10}, essentially equal to the value for B stars, and consistent with the ISM gas phase abundance of \\citet{Meyer:1997}. Thus N is not significantly depleted onto dust grains along this line of sight.  Further discussion of the implications of these results and comparisons with other sightlines can be found in the companion paper by \\citet{Moos:2001}."
    },
    "0201/astro-ph0201121_arXiv.txt": {
        "abstract": "We investigate the pattern of anomalies in the light curves of caustic-crossing binary microlensing events induced by spot(s) on the lensed source star.  For this purpose, we perform simulations of events with various models of spots.  From these simulations, we find that the spot-induced anomalies take various forms depending on the physical state of spots, which is characterized by the surface brightness contrast, the size, the number, the umbra/penumbra structure, the shape, and the orientation with respect to the sweeping caustic.  We also examine the feasibility of distinguishing the two possibly degenerate types of anomalies caused by a spot and a transiting planet and find that the degeneracy in many cases can be broken from the characteristic multiple deviation feature in the spot-induced anomaly pattern caused by the multiplicity of spots. ",
        "introduction": "Microlensing experiments (MACHO: Alcock et al.\\ 1993; EROS: Aubourg 1993) were originally proposed to search for Galactic dark matter in the form of massive compact halo objects by monitoring lensing-induced light variations of stars located in the Large Magellanic Cloud (Paczy\\'nski 1986).  Besides this original goal, it was demonstrated that microlensing can also be applied to other fields of astronomy, especially stellar astrophysics (see the recent review of Gould 2001).  Over the last decade, this aspect of microlensing achieved important progress thanks partially to the theoretical studies on various methods to extract additional information about lensed source stars and more importantly to the large number of event detections from additional experiments directed towards the Galactic bulge (OGLE: Udalski et al.\\ 1993; DUO: Alard \\& Guibert 1997) and detailed light curves of events obtained from intensive followup observations (PLANET: Albrow et al.\\ 1998; GMAN: Alcock et al.\\ 1997; MPS: Rhie et al.\\ 1999; MOA: Bond et al.\\ 2001).   One of the lensing applications to stellar astrophysics is the detection and characterization of stellar spots.  Spot detection via microlensing is possible for high magnification events produced by the source's crossing of the lens caustic, which refers to the source position on which the magnification of a point source event becomes infinity.  For these events, one can resolve the source star surface because different parts of the source is magnified by different amounts due to the large gradient of magnification over the source during the caustic crossing (Gould 1994; Nemiroff \\& Wickramasinghe 1994; Witt \\& Mao 1994).  For a single lens, the location of the caustic is that of the lens itself (point caustic).  For a binary lens, the set of caustics forms closed curves in which each curve is composed of three or more concave line segments (fold caustic) that meet at cusps. Heyrovsk\\'y \\& Sasselov (2000) first investigated the possibility of lensing spot detections and showed that, for point-caustic-crossing single lens events, spots can cause fractional deviations in magnification larger than 2 \\%, which are detectable from followup observations.  Han et al.\\ (2000) further investigated the possibility of detecting stellar spots from the observations of fold-caustic-crossing binary lens events and showed that the fractional deviations are comparable to those of point-caustic-crossing events.  However, these works were concentrated only on the possibility of spot detections, and thus no detailed investigation about the various forms of spot-induced anomalies in lensing light curves has been done.  Bryce \\& Hendry (2000) mentioned some of the variations in their unpublished paper, but these analysis was only for point-caustic-crossing events, which are much less common than fold-caustic-crossing events.  In addition, the previous studies are based on very simplified assumptions  of spots and their host stars, e.g.\\ a circular spot with a uniform surface brightness on also a uniform background stellar surface.  Therefore, extensive study about the microlensing signature of stellar spots for caustic-crossing binary lens events based on realistic models of spots and host stars is required.   Another reason for the necessity of studying spot-induced light curve deviations was recently raised by Lewis (2001).  He pointed out that if a source star possessing a planet is microlensed when the planet is transiting the source star surface, the resulting light curve will have similar deformation to those induced by a spot.  Then, unless one can distinguish the two types of deviations, analyses based on one naively adopted assumption for the cause of the deviation might result in false information about the source star.   In this paper, by performing realistic simulations of microlensing light curves of events occurred on spotted source stars, we investigate how the pattern of the spot-induced anomalies in lensing light curves varies for different states of spots.  Based on the results of this investigation, we also examine the feasibility of distinguishing the two types of anomalies caused by a spot and a transiting planet.   The paper is organized as follows.  In \\S\\ 2, we describe the basics of microlensing, which are  required to describe the lensing behavior of caustic-crossing binary lens events occurred on spotted source stars. In \\S\\ 3, by simulating lensing events occurred on source stars with various states of spots, we investigate the dependencies of the spot-induced anomaly pattern on various spot parameters, which characterize the physical state of spots.  In \\S\\ 4, we discuss the features of the spot-induced anomalies that can be used to distinguish the spot-induced anomalies from those caused by a transiting planet.  We conclude in \\S\\ 5. ",
        "conclusions": "We have investigated the patterns of spot-induced anomalies in caustic-crossing binary lens events.  For this purpose, we performed simulations of events with various models of the physical state of spots which is characterized by the surface brightness contrast, the size, the number, the umbra/penumbra structure, the shape, and the orientation with respect to the sweeping caustic.  From these simulations, we learned that the spot-induced anomalies take various forms depending on these factors. We also examined the feasibility of distinguishing the two possibly degenerate types of anomalies caused by a spot and a transiting planet and found that the degeneracy in many cases can be broken from the characteristic multiple deviation feature in the spot-induced anomaly pattern caused by the multiplicity of spots.    \\bigskip This work was supported by a grant (2001-DS0074) from Korea Research Foundation (KRF)."
    },
    "0201/astro-ph0201317_arXiv.txt": {
        "abstract": "A (sub-)millimeter line and continuum study of the class I protostar Elias 29 in the $\\rho$ Ophiuchi molecular cloud is presented, whose goals are to understand the nature of this source, and to locate the ices that are abundantly present along this line of sight. Within 15--60$''$ beams, several different components contribute to the line emission. Two different foreground clouds are detected, an envelope/disk system and a dense ridge of \\hcop--rich material.  The latter two components are spatially separated in millimeter interferometer maps.  We analyze the envelope/disk system by using inside-out collapse and flared disk models.  The disk is in a relatively face-on orientation ($\\rm < 60^o$), which explains many of the remarkable observational features of Elias 29, such as its flat SED, its brightness in the near infrared, the extended components found in speckle interferometry observations, and its high velocity molecular outflow. It cannot account for the ices seen along the line of sight, however. A small fraction of the ices is present in a (remnant) envelope of mass 0.12--0.33 $M_{\\odot}$, but most of the ices ($\\sim$70\\%) are present in cool ($T<$40 K) quiescent foreground clouds.  This explains the observed absence of thermally processed ices (crystallized H$_2$O) toward Elias 29. Nevertheless, the temperatures could be sufficiently high to account for the low abundance of apolar (CO, N$_2$, O$_2$) ices.  This work shows that it is crucial to obtain spectrally and spatially resolved information from single-dish and interferometric molecular gas observations in order to determine the nature of protostars and to interpret infrared ISO satellite observations of ices and silicates along a pencil beam. ",
        "introduction": "~\\label{se29:intro}  A rich chemical and physical interplay exists between gas and grains in which molecules are formed on grains, creating ice mantles that are preserved in environments ranging from quiescent dense molecular clouds to envelopes and disks around protostars.  Various processes, among which are bombardment by cosmic rays, ultraviolet irradiation, heating, and shocks, can physically or chemically alter the icy mantles, or return molecules into the gas phase (see \\citealt{tiel97, dish98}, and references therein).  A study of the chemical evolution of dense clouds to planet forming disks would ideally involve observations of molecular gas and ices in a range of environments, from quiescent clouds to disk dominated protostars.  The most pristine, initial conditions are presumably well sampled by field stars behind clouds tracing quiescent molecular cloud material (e.g. \\citealt{whit98}). Lines of sight to protostars are more difficult to characterize, however, since they may trace quiescent foreground material, in addition to the gas and ices in their envelopes and disks (e.g. \\citealt{boog00b}). It is thus crucial to characterize the line of sight conditions in order to locate the ices and derive physical conditions and eventually the evolution of the molecular gas and solid state in the interstellar medium.  \\begin{table*}[t!] \\center {\\footnotesize \\caption{Observational Summary}~\\label{t:obssum} \\begin{tabular}{lllll} \\tableline \\noalign{\\smallskip} Telescope    & Beam $\\varnothing$ (\\arcsec) & $\\lambda$ or $\\nu$      & Species$^a$\t   & Date or Reference \\\\ \\noalign{\\smallskip} \\tableline \\noalign{\\smallskip} NRAO-12m     & 43--86\t      & 70--145 GHz\t    & CO, CS, \\hcop        & 05/1995 \\\\ &\t\t      &\t    \t\t    & \\formalde, \\methanol &\t   \\\\ JCMT\t     & 14--22\t      & 200--400 GHz\t    & CO, CS, \\hcop        & 1995-1997 \\\\ &\t\t      &\t    \t\t    & \\formalde, \\methanol &\t   \\\\ JCMT\t     & 7\t      &\t692 GHz\t    \t    & CO 6-5 \t\t   & Ceccarelli et al., in prep. \\\\ CSO\t     & 21--35\t      &\t200--400 GHz\t    & \\formalde, \\methanol & 1999-2001 \\\\ &\t\t      &\t    \t\t    & CO, \\hcop\\ maps      &\t       \\\\ OVRO\t     & 4$\\times$8, 3$\\times$6 & 87, 110 GHz         & CO, HCO$^+$, SiO     & 09/1999-07/2000 \\\\ OVRO\t     & 4$\\times$8, 3$\\times$6 & 2.7, 3.3 mm         & Continuum\t      \t   & 09/1999-07/2000 \\\\ IRAM-30m     & 15\t      & 1.3 mm\t\t    & Continuum\t\t   & \\citealt{mott98} \\\\ ISO SWS      & 14--33\t      & 2.3--45 \\mum\\\t    & Continuum     \t   & \\citealt{boog00b} \\\\ &                &                \t    & Ices, Silicates\t   &                   \\\\ ISO LWS\t     & 80             & 45--200 \\mum\\\t    & Continuum\t   \t   & \\citealt{boog00b} \\\\ \\noalign{\\smallskip} \\tableline \\noalign{\\smallskip} \\multicolumn{5}{l}{$\\rm ^a$ main species and/or isotopes measured. For more details see Table~\\ref{t:sdish}.}\\\\ \\end{tabular} } \\end{table*}  In this paper, we study the line of sight of the class I protostar Elias 29 in the $\\rho$ Oph molecular cloud, using (sub)millimeter single dish and interferometer gas phase observations.  This object is one of the most luminous protostars ($\\sim 36~L_{\\odot}$; \\citealt{chen95}) in the nearby $\\rho$ Oph complex ($d\\sim 160$ pc; \\citealt{whit74}), yet little is known about its nature and line of sight conditions.  Abundant ice has been detected in its direction \\citep{zinn85, tana90}. A detailed analysis of ice band profiles indicates that the ices are not strongly thermally processed (i.e. the ices are not crystallized or segregated), despite the presence of abundant warm molecular gas toward the object \\citep{boog00b}.  This contrasts with high mass, luminous ($\\rm >10^4~L_{\\odot}$) protostars, where significant thermal processing of the ices accompanies the presence of abundant warm molecular gas \\citep{boog00a, tak00}.  The result obtained for Elias 29 can only be understood once the location of the ices and the physical conditions of the various gas components are known.  Therefore, in this paper we try to identify any foreground material, the presence of a circumstellar envelope as well as the presence and orientation of a circumstellar disk, and the column density of each component.  We will then address the question where the ices are located, and what their relation is to the young star. This study will, as a consequence, reveal important information on the nature and evolutionary stage of Elias 29, which has many interesting and unique properties \\citep{elia78}.  Details of the single dish and interferometer observations are presented in \\S 2, and the maps and spectra are decomposed and interpreted in \\S 3.  The physical conditions are determined for the different components along the line of sight, among which are two foreground clouds (\\S 4.1), a remnant envelope and face-on disk (\\S 4.2), as well as a dense ridge from which Elias 29 probably formed (\\S 4.3). The gas phase conditions and abundances are linked to the ice observations in \\S 4.4. The depletion of gas phase species is compared to young class 0 objects and quiescent clouds in \\S 4.5. The results are summarized in \\S 5. ",
        "conclusions": "We have analyzed infrared and millimeter wave line and continuum observations to construct a model of the class I protostar Elias 29 and its environment. This model has to contain a number of different components (summarized in Fig.~\\ref{f:schem}): a disk to account for the $\\sim$2--50 \\mum\\ SED, an envelope contributing to the emission at 1.3 mm (in particular its size), a dense ridge from which Elias 29 may have condensed, and foreground material which provides most of the extinction. Elias 29 can then be well described by a 500 AU radius face-on flared disk with a mass of 0.012 $M_\\odot$, embedded in a 6000 AU radius, 0.12 $M_\\odot$ envelope.  This large disk provides the simplest explanation for the observed flat SED, weakly detected 3 mm continuum emission, and 400 AU radius 5 \\mum\\ thermal continuum emission. The present data does however not fully exclude models with smaller disks. The minimum possible disk has a 30 AU radius and a mass of 0.002 $M_\\odot$ surrounded by an envelope of 0.33 $M_\\odot$.  In this case the combination of disk and envelope emission produces a flat SED.  The entire system is embedded in a long, dense, cold, and \\hcop--rich ridge. Elias~29 is slightly offset from the crest of this ridge.  In front of the disk, envelope, and ridge system are two foreground clouds at a few \\kms\\ lower radial velocities that cover the entire field of view. The large column of the foreground clouds, corresponding to $A_{\\rm V}\\sim$11, may be responsible for the `class I' appearance of Elias~29, which would otherwise appear as a T~Tauri or Herbig Ae/Be star (i.e., optically visible). These same foreground clouds are also the most likely repository of most of the ices seen along the line of sight ($\\sim 70$\\%).  The low temperature of the foreground clouds explains the observed absence of crystallized ices, and the presence of large abundances of polar, H$_2$O--rich ices, i.e. thermal processing did not play a major role for the ices toward Elias 29.  The foreground cloud temperature (25$\\pm$15 K) could, however, be high enough to explain the low abundance of apolar, volatile CO--rich ices, presumably due to the proximity of a number of luminous B type stars.  The important question as to whether the ices in the disk or envelope have experienced thermal processing, as is seen in the envelopes of massive objects \\citep{boog00a, tak00}, cannot be addressed given the large column of foreground material and the face-on nature of the system.  This work shows the value of spectrally and spatially resolved information offered by single-dish and interferometric molecular gas observations in interpreting infrared ISO satellite observations of ices along a pencil beam. It shows that, at least for the $\\rho$ Oph cloud, it is crucial to disentangle the different physical components along the line of sight, since otherwise incorrect conclusions may be derived, for example, on the origin and evolution of interstellar and circumstellar ices.  Follow up observations with sensitive submillimeter interferometers (e.g. ALMA) or high frequency ($>400$ GHz) single dish telescopes with small beams will clearly provide essential information on the structure, abundances, and depletion factor of species in the disk and outflow of Elias 29, and its relationship to the envelope and ridge. In turn, such studies would provide invaluable information on the initial conditions of planet formation."
    },
    "0201/hep-ph0201287_arXiv.txt": {
        "abstract": "We study three-flavor neutrino oscillations in the early universe in the presence of neutrino chemical potentials. We take into account all effects from the background medium, i.e.~collisional damping, the refractive effects from charged leptons, and in particular neutrino self-interactions that synchronize the neutrino oscillations.  We find that effective flavor equilibrium between all active neutrino species is established well before the big-bang nucleosynthesis (BBN) epoch if the neutrino oscillation parameters are in the range indicated by the atmospheric neutrino data and by the large mixing angle (LMA) MSW solution of the solar neutrino problem. For the other solutions of the solar neutrino problem, partial flavor equilibrium may be achieved if the angle $\\theta_{13}$ is close to the experimental limit $\\tan^2\\theta_{13}\\lsim 0.065$. In the LMA case, the BBN limit on the $\\nu_e$ degeneracy parameter, $|\\xi_\\nu|\\lsim 0.07$, now applies to all flavors. Therefore, a putative extra cosmic radiation contribution from degenerate neutrinos is limited to such low values that it is neither observable in the large-scale structure of the universe nor in the anisotropies of the cosmic microwave background radiation. Existing limits and possible future measurements, for example in KATRIN, of the absolute neutrino mass scale will provide unambiguous information on the cosmic neutrino mass density, essentially free of the uncertainty of the neutrino chemical potentials. ",
        "introduction": "\\label{sec:introduction}  The cosmic matter and radiation inventory is known with ever increasing precision, but many important questions remain open.  The cosmic neutrino background is a generic feature of the standard hot big bang model, and its presence is indirectly established by the accurate agreement between the calculated and observed primordial light-element abundances. However, the exact neutrino number density is not known as it depends on the unknown chemical potentials for the three flavors. (In addition there could be a population of sterile neutrinos, a hypothesis that we will not discuss here.)  The standard assumption is that the asymmetry between neutrinos and anti-neutrinos is of order the baryon asymmetry $\\eta \\equiv (n_B-n_{\\bar{B}})/n_\\gamma \\simeq 6\\times 10^{-10}$. This would be the case, for example, if $B-L=0$ where $B$ and $L$ are the cosmic baryon and lepton asymmetries, respectively. While $B-L=0$ is motivated by scenarios where the baryon asymmetry is obtained via leptogenesis~\\cite{Buchmuller:2000as}, there are models for producing large $L$ and small $B$ \\cite{Harvey:1981cu,DK,Casas:1999gx,Dolgov:1991fr}.  In order to quantify a putative neutrino asymmetry we assume that well before thermal neutrino decoupling a given flavor is characterized by a Fermi-Dirac distribution with a chemical potential $\\mu_\\nu$, $f_\\nu(p,T) = \\left [\\exp \\left(E_p/T -\\xi_\\nu \\right)+1 \\right]^{-1}$, where $\\xi_\\nu \\equiv \\mu_\\nu/T$ is the degeneracy parameter and $E_p\\simeq p$ since we may neglect small neutrino mass effects on the distribution function. For anti-neutrinos the distribution function is given by $\\xi_{\\bar{\\nu}}=-\\xi_\\nu$.  A neutrino chemical potential modifies the outcome of primordial nucleosynthesis in two different ways~\\cite{Sarkar:1996dd}.  The first effect appears only in the electron sector because electron neutrinos participate in the beta processes which determine the primordial neutron-to-proton ratio so that $n/p\\propto \\exp(-\\xi_e)$.\\footnote{We use the notation $\\xi_e\\equiv \\xi_{\\nu_e}$ etc.\\ to avoid double subscripts. We never discuss charged-lepton chemical potentials so that there should be no confusion.}  Therefore, a positive $\\xi_e$ decreases $Y_{\\rm p}$, the primordial $^4$He mass fraction, while a negative $\\xi_e$ increases it, leading to an allowed range \\begin{equation} -0.01<\\xi_e<0.07~, \\label{stronglimit} \\end{equation} compatible with $\\xi_e=0$ (see Refs.~\\cite{Kang:1992xa,Esposito:2000hh,Esposito:2001sv,DiBari:2001ua} and Sec.~\\ref{sec:newlimits}).  A second effect is an increase of the neutrino energy density for any non-zero $\\xi$ which in turn increases the expansion rate of the universe, thus enhancing $Y_{\\rm p}$. This applies to all flavors so that the effect of chemical potentials in the $\\nu_\\mu$ or $\\nu_\\tau$ sector can be compensated by a positive $\\xi_e$. Altogether the big-bang nucleosynthesis (BBN) limits on the neutrino chemical potentials are thus not very restrictive.  Another consequence of the extra radiation density in degenerate neutrinos is that it postpones the epoch of matter-radiation equality. In the cosmic microwave background radiation (CMBR) it boosts the amplitude of the first acoustic peak of the angular power spectrum and shifts all peaks to smaller scales. Moreover, the power spectrum of density fluctuations on small scales is suppressed~\\cite{Lesgourgues:1999wu,Hu:1999tk}, leading to observable effects in the cosmic large-scale structure~(LSS).  A recent analysis of the combined effect of a non-zero neutrino asymmetry on BBN and CMBR/LSS yields the allowed regions~\\cite{Hansen:2001hi} \\begin{equation} -0.01 < \\xi_e < 0.22, \\qquad |\\xi_{\\mu,\\tau}| < 2.6, \\label{nooscbounds} \\end{equation} in agreement with similar bounds in~\\cite{Hannestad:2001hn,Kneller:2001cd}. These limits allow for a very significant radiation contribution of degenerate neutrinos, leading many authors to discuss the implications of a large neutrino asymmetry in different physical situations. These include the explanation of the former discrepancy between the BBN and CMBR results on the baryon asymmetry \\cite{Lesgourgues:2000eq} or the origin of the cosmic rays with energies in excess of the Greisen-Zatsepin-Kuzmin cutoff~\\cite{Gelmini:1999qa}. In addition, if the present relic neutrino background is strongly degenerate, it would enhance the contribution of massive neutrinos to the total energy density~\\cite{Pal:1999jv,Lesgourgues:2000ej} and affect the flavor oscillations of the high-energy neutrinos~\\cite{Lunardini:2001fy} which are thought to be produced in the astrophysical accelerators of high-energy cosmic rays.  The limits in Eq.~(\\ref{nooscbounds}) ignore neutrino flavor oscillations, an assumption which is no longer justified in view of the experimental signatures for neutrino oscillations by solar and atmospheric neutrinos. For zero initial neutrino chemical potentials, the flavor neutrinos have the same spectra so that oscillations produce no effect. This is true up to a small spectral distortion caused by the heating of neutrinos from $e^+e^-$ annihilations, an effect which is different for electron and muon/tau neutrinos and which causes a small relative change in the final production of $^4$He of order $10^{-3}$~\\cite{Dolgov:1997mb}. This relative change is slightly enhanced by neutrino flavor oscillations~\\cite{Langacker:1987jv,Hannestad:2001iy}. In the presence of neutrino asymmetries, flavor oscillations equalize the neutrino chemical potentials if there is enough time for this relaxation process to be effective~\\cite{Savage:1991by}. If flavor equilibrium is reached before BBN, then the restrictive limits on $\\xi_e$ in Eq.~(\\ref{stronglimit}) will apply to all flavors, in turn implying that the cosmic neutrino radiation density is close to its standard value.  As a consequence, it is no longer necessary to use the neutrino radiation density as a fit parameter for CMBR/LSS analyses, unless one considers exotic models with decaying massive particles.  The effects of flavor oscillations on possible neutrino degeneracies have been considered in~\\cite{Lunardini:2001fy}, where it was concluded that flavor equilibrium was achieved before the BBN epoch if the solar neutrino problem was explained by the large-mixing angle (LMA) solution. The LMA solution is favored by the current solar neutrino data. Thus, it was concluded that in the LMA case a large cosmic neutrino degeneracy was no longer allowed.  We revisit this problem because the flavor evolution of the neutrino ensemble is more subtle than previously envisaged if medium effects are systematically included.  Contrary to the treatment of Ref.~\\cite{Lunardini:2001fy}, the refractive effect of charged leptons can not be ignored, and actually is one of the dominant effects. While the background neutrinos produce an even larger refractive term, its effect is to synchronize the neutrino oscillations~\\cite{Samuel:1993uw,Pastor:2001iu} which remain sensitive to the charged-lepton contribution. Still, equilibrium is essentially, but not completely, achieved in the LMA case so that our final conclusion is qualitatively similar to that of Ref.~\\cite{Lunardini:2001fy}.  One counter-intuitive subtlety is that the neutrino self-potential actually can suppress oscillations in a situation where the excess of neutrinos in one flavor is exactly matched by an excess of anti-neutrinos in another flavor. In this case the synchronized oscillation frequency is zero so that oscillations begin only once the cosmic expansion has diluted the self-term. Therefore, one can construct cases where flavor equilibrium is not achieved before BBN even in the LMA case.  However, this is only possible for specially chosen initial conditions where $|\\xi_\\nu|$ is equal for all flavors, but the absolute signs may be different. For the purpose of deriving limits on $|\\xi_\\nu|$, however, this case is equivalent to the one where equilibrium is achieved, the only important point being that $|\\xi_\\nu|$ is approximately equal for all flavors at the BBN epoch.  Another subtlety appears in a full three-flavor analysis. We show that achieving equilibrium in the LMA case does not depend on the value of the mixing angle $\\theta_{13}$, to which strict limits from reactor experiments apply. Moreover, for the non-LMA solutions of the solar neutrino problem, partial flavor equilibrium may be reached if the angle $\\theta_{13}$ is small but non-zero.  In Sec.~\\ref{sec:neutrinoflavor} we set up the formalism to study primordial flavor oscillations. We then turn in Sec.~\\ref{sec:twoflavor} to the primordial flavor evolution of a simplified system where $\\nu_e$ mixes maximally with one other flavor. This two-flavor case will illustrate many of the important subtleties of our problem.  Then we turn in Sec.~\\ref{sec:threeflavor} to realistic three-flavor situations which involve yet further complications.  In Sec.~\\ref{sec:newlimits} we finally derive new limits on the degeneracy parameters and summarize our findings. ",
        "conclusions": ""
    },
    "0201/astro-ph0201303_arXiv.txt": {
        "abstract": "{This is the first paper in a series of three where we present the first comprehensive inventory of solid state emission bands observed in a sample of 17 oxygen-rich circumstellar dust shells surrounding evolved stars. The data were taken with the Short and Long Wavelength Spectrographs on board of the Infrared Space Observatory~(ISO) and cover the 2.4 to 195 $\\mu$m wavelength range. The spectra show the presence of broad 10 and 18 $\\mu$m bands that can be attributed to amorphous silicates. In addition, at least 49 narrow bands are found whose position and width indicate they can be attributed to crystalline silicates. Almost all of these bands were not known before ISO. The incredible richness of the crystalline silicate spectra observed by ISO allows detailed studies of the mineralogy of these dust shells, and is a telltale about the origin and evolution of the dust. We have measured the peak positions, widths and strengths of the individual, continuum subtracted bands. Based on these measurements, we were able to order the spectra in sequence of decreasing crystalline silicate band strength. We found that the strength of the emission bands correlates with the geometry of the circumstellar shell, as derived from direct imaging or inferred from the shape of the spectral energy distribution. This naturally divides the sample into objects that show a disk-like geometry (strong crystalline silicate bands), and objects whose dust shell is characteristic of an outflow (weak crystalline silicate bands). All stars with the 33.6 $\\mu$m forsterite band stronger than 20 percent over continuum are disk sources. We define spectral regions (called complexes) where a concentration of emission bands is evident, at 10, 18, 23, 28, 33, 40 and 60 $\\mu$m. We derive average shapes for these complexes and compare these to the individual band shapes of the programme stars. In an Appendix, we provide detailed comments on the measured band positions and strengths of individual sources.} ",
        "introduction": "Red giants and supergiants are characterized by low surface temperatures resulting in the presence of many different molecules in their atmosphere. These objects are also known to have dense stellar winds, presumably driven by a combination of pulsations and radiation pressure on the dust which forms in the cooling outflow. Since dust efficiently absorbes radiation at short wavelengths, the central star can easily become obscured and most of the luminosity of the star is re-radiated at mid-IR wavelengths. At these wavelengths, the most important ro-vibrational bands of abundant molecules can be found, and indeed many solid state bands from various dust components have been found using infrared spectrographs. The Infrared Space Observatory (ISO, Kessler et al. 1996) has allowed for the first time a comprehensive inventory of solid state bands in astrophysical objects with uninterrupted wavelength coverage from 2.4 to 200 $\\mu$m and with a spectral resolution which is well suited for the detection of solid state bands.  We have undertaken detailed studies of the dust emission and absorption spectra of evolved stars, ranging from Asymptotic Giant Branch~(AGB) stars to Planetary Nebulae~(PNe) and to (post) Red Supergiants~(RSG). Preliminary results of these studies have been published elsewhere and can be summarized as follows: Many oxygen-rich evolved stars have a surprisingly rich spectrum of solid state emission bands between 10 and 100 $\\mu$m, dominated by both amorphous and crystalline silicates (e.g. Waters et al. 1996; Molster et al. 1999a; 1999b). The crystalline silicates were not known to be present in dust shells around evolved stars before ISO was launched, and allow for the first time a mineralogical analysis of the dust composition around these objects. Some objects show a very high abundance of crystalline silicates (e.g. Molster et al. 2001a), which seems to be related to the geometry of the circumstellar dust shell (Molster et al 1999a). Surprisingly, stars which were believed to have a carbon-rich dust chemistry also showed the presence of crystalline silicate emission, pointing to a complex chemical composition of the circumstellar environment, possibly due to rapid changes in the chemistry of the stellar photosphere (e.g. Waters et al. 1998; Cohen et al. 1999, Molster et al. 2001a). These observations demonstrate the use of the crystalline silicate bands for a better understanding of the evolution of late type giants and supergiants.   \\begin{table*}[b] \\caption[]{The programme stars and logbook of the observations used for this study. The boldface revolution numbers are those used for the final spectrum.} {\\footnotesize \\begin{tabular}{|l|l|l|l|l|l|} \\hline Object & Type &Revolution     & AOT & T$_{\\rm {int}}$ & Comment\\\\ &      &               & (AOT nr)& (sec)       & \\\\ \\hline IRAS09425-6040 & C-rich AGB star with O-rich dust   & \\bf 084   & SWS01 & 1816  & \\\\ &                                    & \\bf 254   & SWS01 & 6538  & \\\\ NGC~6537        & Planetary Nebula, hot central star &  470      & SWS01 & 1912  & mispointed\\\\ &                                    &  470      & LWS01 & 1318  & \\\\ &                                    & \\bf 703   & SWS01 & 3454  & \\\\ &                                    &\\bf 703    & LWS01 & 2230  & \\\\ NGC~6302        & Planetary Nebula, hot central star & \\bf 094   & SWS01 & 6528  & \\\\ &                                    & 479       & SWS06 & 8532  & $2.4 - 7.0$ and $12.0 - 27.5 \\mu$m\\\\ &                                    &   479     & SWS06 & 12165 & $7.0 - 12.0$ and $27.5 - 45.2 \\mu$m\\\\ &                                    & \\bf289-678& LWS01 & 13243 & combination of rev 289, 482,\\\\ &                                    &           &       &       & 489, 503, 510, 671, and 678\\\\ MWC~922         & Peculiar object                    & \\bf 153   & SWS01 & 1834  & combined with rev 703\\\\ &                                    & \\bf 478   & LWS01 & 1316  & \\\\ &                                    & \\bf 703   & SWS01 & 1912  & combined with rev 153\\\\ AC Her         & Binary post-AGB star               & 106       & SWS01 & 1834  & \\\\ &                                    & \\bf 471   & LWS01 & 1910  & \\\\ &                                    & \\bf 520   & SWS01 & 6538  & \\\\ HD~45677        & B[e] star, nature unclear          & \\bf 711   & SWS01 & 6538  & \\\\ 89~Her         & Binary post-AGB star               & 082       & SWS01 & 1044  & \\\\ &                                    & \\bf 336   & LWS01 & 1860  & \\\\ &                                    & \\bf 518   & SWS01 & 6538  & \\\\ MWC~300         & Evolved star, B supergiant         & \\bf 516   & SWS01 & 3454  & \\\\ Vy~2-2         & Proto-planetary nebula             & \\bf 320   & SWS01 & 1140  & \\\\ &                                    & \\bf 547   & LWS01 & 1318  & $\\approx 7$ arcsec mispointed \\\\ HD~44179        & Red Rectangle; binary post-AGB star& 702       & SWS06 & 1174  & $31.4 - 35.1 \\mu$m \\\\ &                                    & 702       & SWS06 & 856   & $19.5 - 25 \\mu$m\\\\ &                                    & \\bf 702   & SWS01 & 6538  & \\\\ &                                    & \\bf 709   & LWS01 & 3428  & \\\\ &                                    & 870       & SWS06 & 8406  & $12 - 19.5 \\mu$m\\\\ HD~161796       & Post-AGB star                      & 071       & SWS01 & 1044  & \\\\ &                                    & \\bf 080   & LWS01 & 1554  & \\\\ &                                    & \\bf 342   & SWS01 & 1912  & \\\\ &                                    & \\bf 521   & SWS06 & 1744  & $29.0 - 45.2 \\mu$m\\\\ OH~26.5+0.6     & OH/IR star, high mass loss rate    & \\bf 330   & SWS01 & 1912  & \\\\ &                                    &   \\bf 330 & LWS01 & 1268  & \\\\ &                                    &     340   & LWS01 &  822  & \\\\ Roberts~22     & post-AGB star, A supergiant        &     084   & SWS01 & 1044  & mispointed    \\\\ &                                    & \\bf 103   & LWS01 & 478   &  \\\\ &                                    & \\bf 254   & SWS01 & 3454  & mispointed    \\\\ HD~179821       & post-AGB or post-RSG star          &     113   & SWS01 & 1834  & \\\\ &                                    & \\bf 520   & SWS01 & 6538  & \\\\ &                                    & \\bf 319   & LWS01 & 1266  & \\\\ AFGL~4106       & post-RSG, binary                   &  060      & SWS01 & 1130  & \\\\ &                                    &     104   & SWS01 & 1834  & \\\\ &                                    & \\bf 104   & LWS01 &  476  & \\\\ &                                    & \\bf 249   & SWS01 & 3454  & \\\\ NML~Cyg        & Red Supergiant, high mass loss rate& \\bf 052   & SWS01 & 6544  & \\\\ &                                    &     342   & SWS01 & 1140  & \\\\ &                                    & \\bf 530   & SWS06 & 1688  & $29.5 - 45.2 \\mu$m\\\\ &                                    & \\bf 555   & LWS01 & 2798  & \\\\ &                                    &     741   & SWS01 & 1140  & \\\\ IRC+10420      & post-RSG, A supergiant             & \\bf 128   & SWS01 & 3462  & \\\\ &                                    & \\bf 316   & SWS06 & 1718  & $29.3-44.7 \\mu$m \\\\ &                                    & \\bf 724   & LWS01 & 3430  & \\\\ \\hline \\end{tabular}} \\label{tab:log} \\end{table*}  \\afterpage{\\clearpage}  This paper is the first in a series of three where we present a detailed and comprehensive overview of the solid state emission bands in oxygen-rich dust shells surrounding evolved stars and related objects. The purpose of these papers is to quantify as best as possible the presence and characteristics of the numerous new emission bands that have been discovered using the ISO data. In paper~II (Molster et al. 2002a) of this series, we describe the average band profiles of seven \"complexes\" that can be recognized in the combined Short- (de Graauw et al. 1996) and Long- (Clegg et al. 1996) wavelength spectrometers (hereafter referred to as SWS and LWS respectively) that were on board of ISO. Based on the strength of the crystalline silicate bands, we divide in paper~II the sample of 17 stars into two groups. This division is also one which separates objects with a highly flattened dust distribution (referred to here as \"disk\" sources) from those with a more spherical distribution of dust (the non-disk or spherical outflow sources, hereafter referred to as \"outflow\" sources). In the present study, we present the 17 programme stars, we give an overview of the individual spectra, and we quantify the position and strength of the bands. In Paper III (Molster et al. 2002b) we will investigate several trends in the spectrum and correlate them with other information available about these sources.  This paper is organized as follows: Sect.~2 presents the sample, the observations and data reduction; Sect.~3 describes the nature of the individual sources, and the shape of the complexes as compared to the mean. Sect~4 summarizes the results of this study. In Appendix~A we present the band strength data and some more detailed discussion about individual spectral features with respect to their reliability and blending. ",
        "conclusions": "We have made an extensive study of the infrared spectra of O-rich environments of evolved stars.  The spectra show a remarkably rich variety of features, most of which never seen before.  We have defined seven complexes at 10, 18, 23, 28, 33, 40 and 60 $\\mu$m.  Each of these complexes is a convolution of several contributions, which vary independently, giving each spectrum a unique appearance.  We have derived average shapes for the complexes. The large sample enabled us also to identify weak structures.  We have quantified the properties of the complexes and of the components that make up these complexes in terms of peak position, band strength and band width.  We find that the band strength correlates with the geometry of the circumstellar environment, and have accordingly conducted a systematic separate inventory of the bands and complexes for disk sources and outflow sources.  We have presented a detailed description of the shapes of the bands and complexes of the individual stars, and how the spectra of individual objects differ from the average complex shapes. In Paper II, we will discuss, compare and identify the average complexes.  The richness of the spectra undoubtedly contains important information related to the conditions in which the grains were formed and their thermal history. The crystalline nature of the carriers of most of the narrow bands (see also papers II and III) allows a detailed chemical analysis of these components, in a way which is unprecedented.  In this way, it may become possible to trace the evolution and processing of grains from their place of birth (in the outflows of evolved stars), through the interstellar medium, into star forming regions and proto-planetary disks.  For instance, Molster et al. (1999a) suggested that in the disk sources low temperature crystallization takes place.  The crystallization of the silicates in outflow sources is likely to have taken place close to the star, i.e. high temperature crystallization.  Therefore the difference found in the spectra and thus dust properties might have to do with the main crystallization process in these sources.  The data set presented in this study will serve as a starting point for more detailed investigations relating to the identification of the carriers of the bands, and for correlation studies.  These will be presented in papers II and III. In Paper~II we will discuss the average complexes and will identify the carriers of the bands using laboratory spectra of a variety of plausible components. In Paper~III we will study the correlations between the wavelengths and strengths of the bands, and derive fits to the bands using simple radiative transfer and laboratory spectra.  \\vspace{1.0cm} \\noindent{\\bf Acknowledgements.}\\\\ FJM wants to acknowledge the support from NWO under grant 781-71-052 and under the Talent fellowship programm. LBFMW acknowledges financial support from an NWO `Pionier' grant."
    },
    "0201/astro-ph0201245_arXiv.txt": {
        "abstract": "{ We present new valuable BVI photometry of ten Seyfert 1 galaxies and narrow band H$\\alpha$ images for six of these objects as well. The results indicate that the distribution of the luminosity of the sample has an amplitude of almost 4 magnitudes with an average of M$_B$=-20.7. The observed morphologies are confined to early type galaxies. A barred structure is found in only 2 objects. Despite that early morphological types are dominant in this sample, integrated (B-V) colors are very blue. For instance, the SO galaxies show, on average, a (B-V)=0.78. This effect seems to be caused by the luminosity contribution of the active nucleus and/or the disk to the total luminosity of the galaxy. In the B band, the contribution of the active galactic nucleus to the total luminosity of the galaxy varies from 3\\% to almost 60\\% and the bulge to disk luminosity ratio (L$_{bulge}$/L$_{disk}$) ranges from 0.6 to 22. Signs of tidal interactions seems to be a common characteristic since they are observed in 6 of the objects and one of them seems to be located in a poor cluster not yet identified in the literature. In contrast, H$\\alpha$ extended emission is rare, with only 1 galaxy showing clear evidence of it. Luminosity profile decomposition shows that the model Gauss + bulge + disk properly reproduces the surface brightness of the galaxies. However, in order to account for the luminosity profile, most of the disk galaxies needs the inner truncated exponential form with a central cutoff radius ranging from 3 to 10 kpc. This is interpreted in terms of reddened regions that are well identified in the B-V color maps. These regions present very similar colors among them, with (B-V)$\\sim$1.2. This fact could be associated to the presence of dust confined in the inner regions of the galaxies. ",
        "introduction": "The relationship between an active galactic nucleus (AGN) and its host galaxy is one of the key issues in the study of nuclear activity. One would expect that certain properties of active galaxies such as mass, luminosity, bulge to disk ratio and colors could influence nuclear activity or vice versa. In fact, the unified model for an AGN requires gas accretion, most probably from the host galaxy, onto a massive and compact object (Malkan 1983). Similarly, the nuclear burst scenario (Terlevich et al. 1992) requires gas fueling together with an efficient way to concentrate that gas into a small region of space, in a relatively short time-scale. Besides this, gas fueling seems to be connected to the presence of a bar, which is needed to get the non-axisymmetric potential that rises up from the theoretical works (Barnes \\& Hernquist 1991). However, recent observational studies based on optical data show that barred galaxies are not an specific signature for nuclear activity (Ho et al. 1997; Hunt et al. 1999). In this sense, Regan \\& Mulchaey (1999) based on the analysis of 12 Seyferts galaxies imaged with the HST propose central spiral dust lanes as an alternative mechanism to drive the gas to the central regions.  Another important question to consider is the role played by the environment as a trigger of the nuclear activity. De Robertis, Hayhoe \\& Yee (1998) have found that AGNs are not more likely to be associated to interactions than normal galaxies. However, Pastoriza, Donzelli \\& Bonatto (1999), studying a sample of interacting galaxies, found that almost 40\\% of the galaxies may host a low luminosity AGN.  It is also worth noting that most of the studies on the topics exposed above are focused on the galaxy nuclei and some of them on the circumnuclear regions, but only a few works have paid attention to the properties of the stellar populations of the hosts galaxies. As an example, Gonz\\'alez-Delgado et al. (1997)(hereafter GD97), presented H$\\alpha$ images of a sample of 55 active galaxies and S\\'anchez-Portal et al. (2000), gave the results for broad band VRI and narrow band H$\\alpha$ photometry for a sample of 24 nearby active galaxies. Also, Hunt et al. (1997) and M\\'arquez et al. (2000) presented near infrared broad band images of a sample of 26 and 18 active galaxies respectively.\\\\ Our goal in this paper is to present a new valuable set of photometric data for a sample of 10 Seyfert 1 galaxies and to describe the main properties of both the stellar component and the gas of the hosts galaxies. The paper is organized as follows: In Section 2 we describe the sample selection. In Section 3 we summarized the observations and data reduction. Section 4 discusses the photometric results and Section 5 describes the particular properties of each galaxy of the sample. Finally, in Section 6 are summarized our conclusions. ",
        "conclusions": "We have presented new photometric BVI data for 10 Seyfert 1 galaxies together with narrow band H$\\alpha$ images for 6 of these objects as well. The absolute B magnitudes of the galaxies, M$_{\\rm B}$, are found to be spread over a large interval, from -18.74 to -22.34. Integrated (B-V) and (V-I) colors as well as morphological types are derived for the first time in most of the objects. We found that the  morphologies  are confined to early type galaxies: one elliptical, five SO, one Sa and three Sb. Overall, 50\\% of the objects can be considered as compacts. Bars are found only in 2 cases (22\\%). The (B-V) colors of the galaxies show to be biased to the blue. In fact, the SO galaxies of the sample shows, on average, (B-V)=0.78, significantly bluer than the average for this morphological type. We interpret this effect as due to a high contribution of the AGN and/or the disk to the total luminosity of the galaxy.  Signs of tidal interactions are detected in six galaxies of the sample. The case for CTS A08.12 is interesting since it seems to be located in a poor cluster not yet identified in the literature. However, it is not possible to confirm if they are physical interacting objects due to the lack of radial velocities of the suggested companions.  Luminosity profiles were adequately fitted to the Gauss + bulge + disk components. In six out of eight disk models it was necessary to truncate the exponential profile in order to improve the fit. The radius of the central cutoff ranged from 3 up to 10 kpc and it usually corresponds to reddened regions, generally well identified in the B-V color maps. These regions present very similar colors among the sample galaxies, (B-V)$\\sim$1.2. We associated them to the presence of dust in the inner few kiloparsecs of the galaxies.  The profile decomposition allowed us to derive the luminosity contribution of the AGN, bulge and disk separately. We found that in the blue band the AGN contribution to the total luminosity varies from 3\\% up to 56\\%. In addition, the bulge to disk ratio ranges from L$_{bulge}$/L$_{disk}$=0.6 to 22.  H$\\alpha$ images show that only 1 out of 6 galaxies presents disk emission. Additional data is needed in order to confirm if this emission is photoionized by the nuclear continuum or any starbutst component."
    },
    "0201/astro-ph0201073_arXiv.txt": {
        "abstract": "Based on the space MSX observation in bands A(8$\\mu$m), C(12$\\mu$m), D(15$\\mu$m) and E(21$\\mu$m), and the ground SiO maser observation of evolved stars by the Nobeyama 45-m telescope in the v=1 and v=2 J=1-0 transitions, the relation between SiO maser emission and mid-IR continuum radiation is analyzed. The relation between SiO maser emission and the IR radiation in the MSX bands A, C, D and E is all clearly correlated. The SiO maser emission can be explained by a radiative pumping mechanism according to its correlation with infrared radiation in the MSX band A. ",
        "introduction": "The relation between SiO maser emission and IR radiation was predicted by the radiative pumping model for the SiO maser by \\citet{deg76} in which SiO maser is pumped by the $\\triangle v=1$ vibrational transition at 8$\\mu$m.  \\citet{buj81} analyzed a radiative pumping model in which the SiO molecules are located in the inner circumstallar envelope and proposed that the inversion is achieved by a direct population transfer from the $v-1$ state via the absorption of stellar 8$\\mu$m photons. This is followed by an optically thick radiative decay back to $v-1$. Their calculation predict that the SiO maser peak intensity be smaller than the 8$\\mu$m flux.  \\citet{buj81} were also the first to look at whether there exists a relation between the SiO maser emission and stellar IR radiation. Their analysis of data from about 20 objects showed that the observations satisfy the requirement to pump SiO maser radiatively, i.e. they lie below the line S$_{\\rm peak}${\\footnotesize (SiO)} $\\leq$ S{\\footnotesize (8$\\mu$m)} \\citep{buj81}.  \\citet{buj87} assembled observational data of more evolved stars and analyzed the relation between SiO maser emission and radiation in some mid-infrared bands and found a good correlation exists between them. However, analysis of the integrated intensity of the bulge SiO maser sources resulted in large scattering between the SiO maser integrated intensity and the infrared fluxes at the IRAS bands 12, 25 and 60$\\mu$m \\citep{jia95}.  The progress in infrared space astronomy and SiO maser surveys makes it possible now to revisit the relation between SiO maser and mid-infrared emission. The Midcourse Space Experiment (MSX) mission surveyed the entire Galactic plane within $|b| \\leq 4.5\\degr$ in five infrared bands B, A, C, D, and E, at 4, 8, 12, 15 and 21$\\mu$m respectively \\citep{pri01}. In addition to its 30 times better spatial resolution than IRAS, the most sensitive band A of MSX is centered at 8$\\mu$m where the proposed infrared radiation pumps the SiO J=1-0 maser. Furthermore, large scale searches for SiO maser emission from evolved stars in different parts of the Galactic plane have been carried out in the J=1-0 rotational transition at the first and second vibrationally excited states  by  the Nobeyama 45-m telescope \\citep{ita01}. This effort detected a large number of faint SiO maser objects with F$_{12}\\sim$ 1Jy (where F$_{\\lambda}$ means the flux intensity in Jy at band $\\lambda$ and $\\lambda$ refers to the IRAS band 12$\\mu$m, 25$\\mu$m, MSX band A, B, C, D, E). ",
        "conclusions": "Based on new database in the mid-infrared provided by the MSX space survey, the relations between the SiO maser peak intensity and mid-infrared radiation in four MSX bands are analyzed. It is found that the SiO maser intensity is correlated with mid-infrared radiation, which confirms the conclusions from previous studies. The result is consistent with the radiative pumping model for the SiO maser emission in evolved stars."
    },
    "0201/astro-ph0201009_arXiv.txt": {
        "abstract": "In this letter we present the first images of the emission of SiO and H$^{13}$CO$^{+}$ in the nucleus of the starburst galaxy M\\,82. Contrary to other molecular species, which mainly trace the distribution of the star forming molecular gas within the disk, the SiO emission extends noticeably out of the galaxy plane. The bulk of the SiO emission is restricted to two major features. The first feature, referred to as the SiO supershell, is an open shell of 150\\,pc diameter, located 120\\,pc west from the galaxy center. The SiO supershell represents the inner front of a molecular shell expanding at $\\sim$40\\,kms$^{-1}$, produced by mass ejection around a supercluster of young stars containing SNR 41.95+57.5. The second feature is a vertical filament, referred to as the SiO chimney, emanating from the disk at 200\\,pc east from the galaxy center. The SiO chimney reaches a 500\\,pc vertical height and it is associated with the most prominent chimney identified in radiocontinuum maps. The kinematics, morphology, and fractional abundances of the SiO gas features in M\\,82 can be explained in the framework of shocked chemistry driven by local episodes of gas ejection from the starburst disk. The SiO emission stands out as a privileged tracer of the disk-halo interface in M\\,82. We speculate that the chimney and the supershell, each injecting  $\\sim$10$^{7}$M$_{\\sun}$ of molecular gas, are two different evolutionary stages in the outflow phenomenon building up the gaseous halo. ",
        "introduction": "The detection of large-scale outflows in over a dozen starburst galaxies has confirmed the overall predictions of the galactic wind model, first proposed by \\citet{che85}. It is widely accepted that the driving mechanism of the outflow phenomenon in starbursts is linked to the creation of expanding shells of hot gas by supernovae. At the end of the starburst cycle, the resulting high supernova rate creates a rarefied wind of hot gas in the disk (with temperatures of $\\sim$10$^7$K) at several thousands of kms$^{-1}$. Different {\\it hot bubbles} could merge and blow out into the halo, entraining surrounding cold gas and dust at several hundreds of  kms$^{-1}$. These bubbles, although initially spherical, may evolve eventually into vertical chimneys of gas (Norman \\& Ikeuchi 1989; Koo \\& McKee 1992; Alton, Davies, \\& Bianchi 1999). The halo outflow would be built up in the end by the local injection of gas from the disk. The details of this secular process, which should drive large-scale shocks in the molecular gas, remain unknown however, partly because we lack of observational constraints.  M\\,82 is the closest galaxy experiencing a massive star forming episode \\citep{rie80, wil99}. Its nuclear starburst, located in the central 1\\,Kpc, has been studied in virtually all wavebands from X-rays to the radio domain. X-rays and optical observations have shown the existence of a large-scale biconical outflow of hot gas coming out of the plane from the nucleus of M82 (Bregman, Schulman, \\& Tomisaka 1995; Shopbell \\& Bland-Hawthorn 1998). This massive outflow is also observed at large scales in the cold gas and dust \\citep{sea01,alt99}. At smaller scales, close to the disk-halo interface of M\\,82, there is observational evidence of local sources of gas injection. Using 1.4 and 5 GHz VLA data, \\citet{wil99} have shown the existence of a large (diameter $\\sim$120 pc) expanding shell of ionized gas, close to the supernova remnant SNR 41.95+57.5. A molecular gas counterpart of this supershell was tentatively identified by \\citet{nei98} and later discussed in \\citet{wei99,wei01} and \\citet{mat00}. \\citet{wil99} detected also the signature of four chimneys of hot gas. The most prominent one, located on the northeastern side, reaches a vertical height of $\\sim$200\\,pc. The molecular gas counterpart of the chimney has remained so far undetected (see \\citet{wei01}).  In this letter we present the first high-resolution ($\\sim$5'') image of the emission of silicon monoxide (SiO) in the nucleus of M\\,82. SiO is known to be a privileged tracer of large-scale shocks in the interstellar medium of galaxies \\citep{mpi97, gbu00, gbm01}. The SiO emission is used in this work to study the occurrence of shocks in the halo-disk interface of M\\,82. We present the first evidence of a SiO expanding supershell, related with the superbubble of hot gas, and the first detection of a molecular gas chimney associated with a violent mass ejection event in M\\,82. ",
        "conclusions": "Numerical simulations studying the evolution of outflows in starbursts \\citep{tom93, suc96} have predicted that the resulting hot galactic wind, observed in X-rays and diffuse H$\\alpha$ emission, may drag the cooler and denser material of the blown-out disk up to 1--2 kpc above the plane of a galaxy. These models show that, near the base of the outflow, at a scale height of $\\sim$500\\,pc, filaments of cold disk material should be present. The detection of SiO emission from a prominent chimney and a giant shell in M\\,82 provides a nice confirmation of these models, proving that the entrained cold material can survive in molecular form in spite of the high-velocities (several hundreds kms$^{-1}$ in the SiO chimney) involved upon ejection. The chemical processing of dust grains by shocks can be at work during the blow-out of the disk; the latter naturally explains the high abundances of SiO measured in the molecular gas of the chimney, and to a lesser extent, in the supershell. Contrary to other species, which mainly trace the distribution of the star forming molecular gas within the disk, SiO stands out as a privileged tracer of the disk-halo interconnection in M\\,82.  Several scenarios can be envisaged to account for the different morphologies and properties characterizing the chimney and the supershell. We discuss the case where the differences might reflect the evolution expected for an ejection event from the disk. Furthermore, the energies required to form the chimney or the supershell may largely differ. Their unlike morphologies may also result from the action of collimating mechanisms shaping differently the outflow of molecular gas.  In the frame of the evolutionary scenario, the SiO chimney, extended up to 500\\,pc above the galaxy plane, would be an evolved ejection episode. In contrast, the supershell (of $\\sim$75\\,pc radius) would be just starting to undergo a blow-out. \\citet{wil99} have come to similar conclusions analyzing the morphologies of the related radiocontinuum features. Additional insight is gained by estimating the kinematical ages of the SiO structures. The reported size ($\\sim$75\\,pc radius) and expansion velocity ($\\sim$40$\\pm$5\\,kms$^{-1}$) allows to infer an age of $\\sim$2\\,10$^{6}$ years for the SiO supershell. This is probably an upper limit, as the gas flow has been likely decelerated during the expansion, as already pointed out by \\citet{wei99}. Deriving the kinematical age of the chimney is less straightforward, however. Although the radial velocities measured along the chimney reveal an ejection-dominated flow, the value of deprojected velocities are strongly model-dependent. A global model for the molecular gas outflow needs a complete high-resolution mapping of the molecular halo in M\\,82. However, we found in our data evidence that the dynamics of the entrained molecular gas is similar to the ionized gas near the location of the SiO chimney. If the model of \\citet{sbh98} (based on a global fit on several optical filaments) held for the SiO chimney, the estimated ejection velocity would be $\\sim$500\\,kms$^{-1}$. The derived kinematical age for the SiO chimney ($\\sim$10$^{6}$ years) would not be significantly different from the one determined for the supershell. As the evolutionary link hypothesis depicted above is probably correct, this result suggests that the time-scale for evolution is shorter in the chimney than in the supershell. This may indicate that the energy required to create the chimney is comparatively larger.   We estimate that the kinetic energy contained in the SiO supershell is E$_{kin}\\sim$2\\,10$^{53}$ergs. Assuming that the typical type-II supernova energy input is 10$^{51}$ergs, and that only $\\sim$10$\\%$ of the explosion energy is transferred  to kinetic energy (see numerical models of \\citet{che74}), the formation of the supershell would require $\\sim$2\\,10$^{3}$ correlated explosions in 10$^{6}$ years. Most notably, the derived kinetic energy of the SiO chimney is much larger: E$_{kin}\\sim$10$^{55}$ergs (taking $\\sim$500\\,kms$^{-1}$ as ejection velocity); this would require 10$^{5}$ correlated supernovae in 10$^{6}$ years, namely a supernova rate of 0.1SN yr$^{-1}$. This result holds even in the improbable scenario where the outflow is nearly parallel to the galaxy plane. In this case we obtain a lower limit of E$_{kin}\\sim$4\\,10$^{53}$ergs for the chimney. This is still a factor 2-3 larger than derived for the supershell. Although the energy requirements to form the SiO chimney might be very stringent if gas velocities are close to hundreds of kms$^{-1}$, the measured supernova rate in the central disk of M\\,82 (0.1SN yr$^{-1}$; see Kronberg, Biermann, \\& Schwab (1985)) may account for it.  We can speculate that the large-scale molecular gas halo of M\\,82 detected by \\citet{sea01}, has been built up by local episodes of gas injection from the disk. The two SiO features reported in this letter, injecting each $\\sim$10$^{7}$M$_{\\sun}$ gas, provides a tantalizing evidence that this secular process is at work in M\\,82. \\citet{sea01} estimated a molecular gas mass of the halo of $\\sim$5\\,10$^{8}$M$_{\\sun}$; the latter implies we would need 20-50 of these local episodes to build up the halo. In view of the available estimates for the age of the starburst (5\\,10$^{7}$-10$^{8}$ years), and the energy deposited by supernovae explosions during this time ($\\sim$10$^{58}$ ergs) we can conclude that both the time-scales and the energy input required to form 20-50 of these ejection episodes are well accounted for by the starburst engine in M\\,82."
    },
    "0201/astro-ph0201523_arXiv.txt": {
        "abstract": "We consider the power of a relativistic jet accelerated by the magnetic field of an accretion disc. It is found that the power extracted from the disc is mainly determined by the field strength and configuration of the field far from the disc. Comparing with the power extracted from a rotating black hole, we find that the jet power extracted from a disc can dominate over that from the rotating black hole. But in some cases, the jet power extracted from a rapidly rotating hole can be more important than that from the disc even if the poloidal field threading the hole is not significantly larger than that threading the inner edge of the disc. The results imply that the radio-loudness of quasars may be governed by its accretion rate which might be regulated by the central black hole mass. It is proposed that the different disc field generation mechanisms might be tested against observations of radio-loud quasars if their black hole masses are available. ",
        "introduction": "The current most favoured models of powering active galactic nuclei (AGNs) involve gas accretion onto a massive black hole, though the details are still unclear. Relativistic jets have been observed in many radio-loud AGNs and are believed to be formed very close to the black holes. The spin energy of a black hole might be extracted to fuel the jet by magnetic fields supported by a surrounding accretion disc (Blandford \\& Znajek 1977; Blandford-Znajek process). This process has widely been believed to be the major mechanism that powers radio jets in AGNs (Begelman, Blandford \\& Rees 1984; Rees et al. 1982; Wilson \\& Colbert 1995; Moderski \\& Sikora 1996). However, Ghosh \\& Abramowicz (1997) doubted the importance of the Blandford-Znajek process. For a black hole of a given mass and angular momentum, the strength of the Blandford-Znajek process depends crucially on the strength of the poloidal field threading the horizon of the hole. They argued that the strength of the field threading a black hole has been overestimated. The magnetic field threading a hole should be maintained by the currents situated in the inner region of a surrounding accretion disc. Livio, Ogilvie \\& Pringle (1999; hereafter L99) re-investigated the problem and pointed out that even the calculations of Ghosh \\& Abramowicz (1997) have overestimated the power of the Blandford-Znajek process. Ghosh \\& Abramowicz (1997) have overestimated the strength of the large-scale field threading the inner region of an accretion disc, and then the power of the Blandford-Znajek process. They argued that the jet power extracted from an accretion disc dominates over the power extracted by the Blandford-Znajek process (L99).  L99 estimated the maximal electromagnetic output from an accretion disc. In this case, the toroidal field component is of the same order of the poloidal field component at the disc surface due to the instability of a predominantly toroidal field (Biskamp 1993; Livio \\& Pringle 1992). Apart from the strength of the field threading a disc, the acceleration of the jet is also governed by the magnetic field configuration and the structure of the disc (Shu 1991; Ogilivie \\& Livio 1998, 2000; Cao \\& Spruit 2001; for relativistic jets near rotating black holes, see Koide, Shibata \\& Kudoh 1999; Koide et al. 2000). It is unclear how much power can actually be extracted from a magnetized accretion disc without considering its field configuration. In this work, we extend L99's work to explore some factors that affect the power of a relativistic jet extracted from an accretion disc. ",
        "conclusions": "Figure 1 shows that the function $f(\\alpha,~\\gamma_{\\rm j})$ varies slowly with $\\gamma_{\\rm j}$ except for $\\gamma_{\\rm}\\rightarrow 1$. The function $f(\\alpha,~\\gamma_{\\rm j})$ is almost a constant while $\\gamma_{\\rm j}>5$. From Eq. (6), it is then found that the power of a relativistic jet accelerated by a magnetized accretion disc is insensitive to its Lorentz factor $\\gamma_{\\rm j}$. This is confirmed by the results plotted in Fig. 5. The jet power varies with $\\gamma_{\\rm j}$ in a range less than 10 \\% for $\\gamma_{\\rm j}>5$. It implies that relativistic jets with different Lorentz factors $\\gamma_{\\rm j}$ accelerated by magnetized discs have almost same jet power, if all other physical parameters are fixed. We can therefore adopt a typical value of $\\gamma_{\\rm j}$ in our calculations, which will not change the main results on jet power, though the terminal velocity of the magnetically driven jet could be given in principle only after the structure of the disc and its field strength and configuration are supplied. This has simplified our present investigation, since we are interested in the power of relativistic jets extracted from discs.  The power of a jet decreases with the self-similar index $\\alpha$ (see Fig. 5). A large $\\alpha$ implies that the poloidal field strength decays rapidly with radius along a field line. From Fig. 4, we see that less mass is loaded into the jet in the case of  a larger $\\alpha$, and the power extracted from the disc is then reduced. The power of the jet is sensitive to the field configuration far from the disc surface.  We find that the mass loss rate in the jet from unit surface area of the disc decreases with radius (see Figs. 2 and 3). Almost all mass carried by the jet is from the region of the disc: $R_{\\rm d}<5~R_{\\rm in}$, i.e., jets are mainly accelerated from the inner region of the disc where most gravitational energy of accretion matter is released. The total mass loss rate in a relativistic jet can be neglected compared with the accretion rate of the disc (usually $\\dot{M}_{\\rm jet}<10^{-2}\\dot{M}_{\\rm acc}$, see Fig. 4). A jet with high mass loss rate results in a low terminal velocity of the jet and wound-up field lines (Figs. 4 and 6), which is consistent with the results of Cao \\& Spruit (1994). As pointed out in Sect. 2, our calculation of the jet power extracted from a disc can reproduce L99's result while $\\alpha=1$ is adopted (see Eq. (8)). In L99's calculation, $B_{\\phi\\rm d}=B_{\\rm pd}$ at the disc surface is assumed, so it is not surprised to find that $B_{\\phi\\rm d}\\simeq B_{\\rm pd}$ in the case of $\\alpha=1$ in our calculations (see Fig. 6). For a large $\\alpha$, the ratio $B_{\\phi\\rm d}/B_{\\rm pd}$ is low (compare the results in Fig. 6 for different values of $\\alpha$).  In our present calculations, the field configuration far from the disc is described by the self-similar index $\\alpha$. In the case of $\\alpha=1$, it is found in Fig. 7 that the power extracted from the disc always dominates over the maximal power of the Blandford-Znajek process, if $\\zeta<3$, i.e., the strength of the poloidal field threading the hole is less than three times of that threading the inner edge of the disc. If $\\zeta=1$ is assumed, i.e., $B_{\\rm ph}=B_{\\rm pd}(R_{\\rm in})$, $L_{\\rm d}$ is an order of magnitude larger than $L_{\\rm BZ}$ even for an extreme Kerr black hole. Thus the conclusion given by L99 that the maximal power extracted from a disc dominates over that from a rapidly rotating black hole is confirmed, while our calculations are not limited to the extreme case $\\alpha=1$. We present the results in Fig. 7 to compare the power outputs by these two processes for different values of $\\alpha$ and $\\zeta$. It is found that the Blandford-Znajek process would be important for a large $\\alpha$ or/and $\\zeta$. From Eqs. (6) and (18), we find that this conclusion is independent of the mechanism for producing disc fields. The conclusion is also independent of the hole mass $m$ and accretion rate $\\dot{m}_{\\rm acc}$, since the hole field strength is only governed by the disc field strength.  If the large-scale field can be produced from the small-scale field created by dynamo processes as done by L99, we find that linear relations are present: $L_{\\rm jet}/L_{\\rm acc}\\propto \\dot{m}_{\\rm acc}$, for a jet accelerated either from a disc or a rotating hole. (see Fig. 8 and Eq. (20)). This relation is quite different from the situation for the field that is assumed to scale with the pressure in the disc, where $L_{\\rm jet}/L_{\\rm acc}\\propto \\dot{m}_{\\rm acc}^{-1}$ is present.  The ratio $L_{\\rm jet}/L_{\\rm acc}$ can be related with an observational quantity of quasars: radio-loudness $R$, which is defined as $R\\equiv f_{\\nu}{\\rm (5GHz)}/f_{\\nu}(\\rm{4400\\AA})$. The fact that the ratio $L_{\\rm jet}/L_{\\rm acc}$ is independent of the black hole mass may imply that the radio-loudness $R$ is not related directly with the black hole mass. The radio-loudness $R$ may probably governed by the accretion rate. The accretion rate is then regulated by the central black hole mass in some way (Yi 1996; Salucci et al. 1999; Haiman \\& Menou 2000; Kauffmann \\& Haehnelt 2000). The relations found between radio-loudness and the central black hole masses in quasars (McLure et al. 1999; Laor 2000; Gu, Cao \\& Jiang 2001; Ho 2002) may reflect intrinsic relations between radio-loudness and accretion rate which is regulated by the black hole mass (Boroson 2002). The different relations between $L_{\\rm jet}/L_{\\rm acc}$ and $\\dot{m}_{\\rm acc}$ are expected for different mechanisms of producing the disc field. The different disc field generation mechanisms can therefore be tested against observations of radio-loud quasars if their black hole masses are available. However, whether the jet power is extracted from a disc or a rotating black hole cannot be tested in this way, since the calculations predict the same relation between $L_{\\rm jet}/L_{\\rm acc}$  and $\\dot{m}_{\\rm acc}$ for these two jet acceleration mechanisms.   In Fig. 9, we compare the relative importance of these two power extract processes in the parameter space ($\\alpha-a$). Only for those cases the strength of the poloidal field of the disc decays slowly with radius along the field line (small $\\alpha$), the power extracted from the disc dominates over that extracted by the Blandford-Znajek process even for a rapidly rotating hole. In the right-upper corner of the figure, i.e., the case with a large $\\alpha$ and $a$, the power extracted from the disc is greater than the power of the Blandford-Znajek process. Otherwise, the power of the Blandford-Znajek process can be neglected compared with that from the disc.  The acceleration of the jet is governed by the structure of the disc and the strength and configuration of the field threading the disc. The physical properties of the jet, such as the power and terminal velocity of the jet, can be determined if all these physical factors of the disc and its field are known. Instead of solving a set of MHD equations describing the disc-jet system, a self-similar index $\\alpha$ is employed to describe the field far above the disc surface in present work. The physics of how the jet acceleration is governed by the detailed structure of the disc near a rapidly rotating black hole has not been included in this work. It would be the main limitation of the present work."
    },
    "0201/astro-ph0201379_arXiv.txt": {
        "abstract": "We present X-ray and optical observations of CXOMP\\,J213945.0-234655, a high redshift ($z=4.93$) quasar discovered through the Chandra Multiwavelength Project (ChaMP).  This object is the most distant X-ray selected quasar published, with a rest-frame X-ray luminosity of L$_{X}=5.9\\times 10^{44}$ erg~s$^{-1}$ (measured in the 0.3--2.5~keV band and corrected for Galactic absorption).  CXOMP\\,J213945.0-234655 is a $g^{\\prime}$ dropout object ($>26.2$), with $r^{\\prime}=22.87$ and $i^{\\prime}=21.36$.  The rest-frame X-ray to optical flux ratio is similar to quasars at lower redshifts and slightly X-ray bright relative to $z>4$ optically-selected quasars observed with Chandra. The ChaMP is beginning to acquire significant numbers of high redshift quasars to investigate the X-ray luminosity function out to $z\\sim5$. ",
        "introduction": "The observed characteristics of known quasars are remarkably similar over a broad range of redshift.  For example, X-ray studies utilizing the ROSAT database \\citep{gr95,ka00}, show little variation of the ratio of X-ray to optical flux for optically selected quasars.  Also, the rest frame UV spectra of quasars, including the broad Ly$\\alpha$, NV and CIV emission lines, are nearly identical for a large range of redshift and present no evidence for subsolar metallicities even up to a $z\\sim6$ \\citep{fa01}.  Even though the individual properties of quasars are similar, the co-moving space density of quasars changes drastically with redshift. At high redshift ($z>4$), a significant dropoff in the co-moving space density of quasars seen in optical (e.g., Schmidt et al. 1995; Warren et al. 1994; Osmer 1982) and radio surveys \\citep{sh96} hints at either the detection of the onset of accretion onto supermassive black holes or a missed high-redshift population, possibly due to obscuration.  X-ray selected quasars from ROSAT have been used to support the latter interpretation based on evidence for constant space densities beyond a redshift of 2 \\citep{mi00}.  Unfortunately, the ROSAT sample size is small with only 8 quasars beyond a redshift of 3.  Significant numbers of quasars with $z>4$ are being amassed to investigate both their intrinsic properties and the evolutionary behavior of the quasar population. The Sloan Digital Sky Survey (SDSS) reports 123 optically selected quasars with $z>4$ \\citep{sc01,an01}.  However, optical surveys suffer from selection effects due to intrinsic obscuration and the intervening Ly$\\alpha$ forest.  Current X-ray surveys with Chandra and XMM do not have a strong selection effect based on redshift and can detect emission up to 10~keV (observed frame) to reveal hidden populations of active galactic nuclei (AGN) including heavily obscured quasars \\citep{no01,st01}.  High-$z$ objects can be detected through a larger intrinsic absorbing column of gas and dust because the observed-frame X-ray bandpass corresponds to higher energy, more penetrating X-rays at the source.\\footnote{The observed-frame, effective absorbing column is $N_{\\rm H}^{\\rm eff}\\sim N_{\\rm H}/(1+z)^{2.6}$ (Wilman \\& Fabian 1999).}  Therefore, optical and X-ray surveys will complement each other, providing a fair census of mass accretion onto black holes at high redshift.  Larger samples of X-ray observations of $z>4$ quasars are needed since there are currently only 24 \\citep{vi01}, of which only 3 are X-ray selected quasars. Chandra and XMM-Newton are beginning to probe faint flux levels for the first time to detect the high-$z$ quasar population.  Initial Chandra and XMM-Newton observations of optically selected quasars have shown a systematically lower X-ray flux relative to the optical at high redshift \\citep{vi01,br01a}.  In this paper, we present the X-ray and optical properties of a newly discovered, X-ray selected $z=4.93$ quasar with the Chandra Observatory.  This quasar is the highest redshift object published\\footnote{A $z\\sim5.2$, X-ray selected quasar detected in the CDF-N was presented at the 199th AAS meeting (Brandt 2001b).} from an X-ray survey.  These results are a component of the Chandra Multiwavelength Project (ChaMP; Wilkes et al. 2001).  A primary aim of the ChaMP is to measure the intrinsic luminosity function of quasars and lower luminosity AGN out to $z\\sim5$. The survey will provide a medium-depth, wide-area sample of serendipitous X-ray sources from archival Chandra fields in Cycles~1 and 2 covering $\\sim 14$ deg$^2$.  The broadband sensitivity between 0.3--8.0 keV enables the selection to be far less affected by absorption than previous optical, UV, or soft X-ray surveys. Chandra's small point spread function ($\\sim$1$\\arcsec$ resolution on-axis) and low background allow sources to be detected to fainter flux levels, while the $\\sim 1^{\\prime\\prime}$ X-ray astrometry greatly facilitates unambiguous optical identification of X-ray counterparts.  The project will effectively bridge the gap between flux limits achieved with the Chandra deep field observations and those of past ROSAT surveys.  Throughout this paper, we assume H$_{\\circ}$=50 km s$^{-1}$ Mpc$^{-1}$ and a flat cosmology with q$_{\\circ}$=0.5. ",
        "conclusions": "We present the discovery of CXOMP\\,J213945.0-234655, at $z=4.93$ the most distant X-ray selected object published to date.  With a measured optical to X-ray flux ratio $\\alpha_{ox}$=1.52, CXOMP\\,J213945.0-234655 is similar to low redshift quasars, in contrast to several optically-selected $z>4$ quasars previously detected by Chandra.  This detection highlights the importance of wide area, intermediate depth surveys like the ChaMP for studies of the high redshift quasar population ($z\\sim$ 3 to 5).  The ChaMP\\footnote{http://hea-www.harvard.edu/CHAMP/} has begun to amass a sample of high redshift, X-ray selected quasars with the goal of measuring the cosmic evolution of accretion-powered sources relatively unhampered by the absorption and reddening that affects optical surveys."
    },
    "0201/astro-ph0201186_arXiv.txt": {
        "abstract": "We present the $2-8$~keV number counts from the 1~Ms {\\it Chandra} observation of the {\\it Chandra} Deep Field North (CDF-N). We combine these with the number counts from a 78~ks exposure of the Hawaii Survey Field SSA22 and with the number counts obtained in independent analyses of the CDF-S and the Hawaii Survey Field SSA13 to determine the number counts from $2\\times 10^{-16}$ to $10^{-13}$~erg~cm$^{-2}$~s$^{-1}$. Over this flux range the contribution to the X-ray background is $1.1\\times 10^{-11}$~erg~cm$^{-2}$~s$^{-1}$~deg$^{-2}$. When the contributions above $10^{-13}$~erg~cm$^{-2}$~s$^{-1}$ from {\\it BeppoSAX} or {\\it ASCA} observations are included, the total rises to $1.4\\times 10^{-11}$~erg~cm$^{-2}$~s$^{-1}$~deg$^{-2}$. However, there appears to be substantial field-to-field variation in the counts in excess of the statistical uncertainties. When the statistical and flux calibration uncertainties (in both the background and source measurements) are taken into account, as much as $0.5\\times 10^{-11}$~erg~cm$^{-2}$~s$^{-1}$~deg$^{-2}$ could still be present in an unresolved component. ",
        "introduction": "\\label{intro}  The X-ray background (XRB) was the first cosmic background detected (\\markcite{giacconi62}Giacconi et al.\\ 1962) and has been extensively characterized (\\markcite{fabian92}Fabian \\& Barcons 1992). Its photon intensity $P(E)$, where $E$ is the photon energy in keV and $P(E)$ has units [photons~cm$^{-2}$~s$^{-1}$~keV$^{-1}$~sr$^{-1}$], can be approximated in the $1-15$~keV range by $P(E)=AE^{-\\Gamma}$ where $\\Gamma\\simeq 1.4$ (e.g., \\markcite{marshall80}Marshall et al.\\ 1980; \\markcite{gendreau95}Gendreau et al.\\ 1995; \\markcite{chen97}Chen, Fabian, \\& Gendreau 1997). However, the XRB normalization is still somewhat uncertain with values in the $2-8$~keV band lying between $1.3\\times 10^{-11}$ (\\markcite{marshall80}Marshall et al.\\ 1980) and $1.8\\times 10^{-11}$~erg~cm$^{-2}$~s$^{-1}$~deg$^{-2}$ (\\markcite{vecchi99}Vecchi et al.\\ 1999).  After the discovery of the XRB there was controversy over whether it arose from a superposition of discrete sources or from thermal bremsstrahlung from a hot intergalactic gas (e.g., \\markcite{field72}Field\\ 1972). We now know that the XRB cannot originate in an uniform hot intergalactic medium since the absence of a strong Compton distortion in the cosmic microwave background spectrum puts a stringent upper limit ($\\sim 10^{-4}$) on such a contribution (\\markcite{wright94}Wright et al. 1994). However, there may be other sources of hard X-ray photons (e.g., \\markcite{abazajian01}Abazajian, Fuller, \\& Tucker 2001) which could contribute to the XRB without producing a Compton distortion, so constraining any residual diffuse background is of considerable interest.  The sources contributing the bulk of the $2-8$~keV XRB were not found prior to the launch of the {\\it Chandra X-ray Observatory}. While the soft ($0.5-2$~keV) background was largely resolved into sources by the {\\it ROSAT} satellite (\\markcite{has98}Hasinger et al.\\ 1998), the spectra were too steep to account for the flat XRB spectrum. With the launch of {\\it Chandra} and the {\\it XMM-Newton Observatory} the situation has evolved rapidly. Early $100-300$~ks {\\it Chandra} observations determined the number counts in the $2-8$~keV band above $10^{-15}$~erg~cm$^{-2}$~s$^{-1}$ (\\markcite{mushotzky00}Mushotzky et al.\\ 2000; \\markcite{giacconi01}Giacconi et al.\\ 2001; \\markcite{tozzi01}Tozzi et al.\\ 2001) and resolved the majority of the $2-8$~keV XRB. \\markcite{has01}Hasinger et al. (2001) have reported parallel results with {\\it XMM-Newton}. Now the two 1~Ms {\\it Chandra} exposures of the CDF-N (\\markcite{brandt01}Brandt et al.\\ 2001, hereafter B01) and the {\\it Chandra} Deep Field South (CDF-S; \\markcite{giacconi02}Giacconi et al.\\ 2002, hereafter G02; \\markcite{campana01}Campana et al.\\ 2001, hereafter C01) have extended the counts to $2\\times 10^{-16}$~erg~cm$^{-2}$~s$^{-1}$. However, there is considerable spread (about 40\\%) in the normalizations, with the counts in the SSA13 (\\markcite{mushotzky00}Mushotzky et al.\\ 2000) and CDF-N (B01) fields being higher than the CDF-S counts (G02; C01) and the {\\it XMM-Newton} Lockman Hole counts (\\markcite{has01}Hasinger et al.\\ 2001).  Here we provide a more detailed analysis of the number counts in the CDF-N, modelling the effects of Eddington bias and incompleteness.  We also analyze a new 78~ks observation of the SSA22 field. We combine our counts with those measured by \\markcite{mushotzky00}Mushotzky et al.\\ (2000) from the 100~ks observation of the SSA13 field and those measured by C01 from the 1~Ms observation of the CDF-S to average over a spread of fields. We are currently analyzing the latter two fields using the same methods described here (L.L. Cowie et al., in preparation), but we note that similar number counts have been obtained for the CDF-S using very different methods (G02; C01).  We break with the tradition of X-ray number counts by using differential and not cumulative counts. Differential counts have many advantages in the statistical independence of the data points and in the ease with which breaks and shape changes may be seen (c.f., \\markcite{jauncey75}Jauncey 1975). ",
        "conclusions": "Fig.~\\ref{fig3} shows the contribution to the XRB versus flux. The peak contribution arises near the break at $1.4 \\times 10^{-14}$~erg~cm$^{-2}$~s$^{-1}$. The total contribution is $1.1\\times 10^{-11}$~erg~cm$^{-2}$~s$^{-1}$~deg$^{-2}$ over the flux range $2\\times 10^{-16}$ to $10^{-13}$~erg~cm$^{-2}$~s$^{-1}$. Extrapolating the power-law fit would add a further $0.2\\times 10^{-11}$~erg~cm$^{-2}$~s$^{-1}$~deg$^{-2}$ above $10^{-13}$~erg~cm$^{-2}$~s$^{-1}$, similar to the value obtained from the {\\it BeppoSAX} number counts. Use of the highest {\\it ASCA} counts would only increase this to $0.3\\times 10^{-11}$~erg~cm$^{-2}$~s$^{-1}$~deg$^{-2}$. Extrapolating the faint end flux counts below $2\\times 10^{-16}$~erg~cm$^{-2}$~s$^{-1}$ would add $0.2\\times 10^{-11}$~erg~cm$^{-2}$~s$^{-1}$~deg$^{-2}$ to the total. For the {\\it HEAO-1} normalization the observed contribution from sources above $2\\times 10^{-16}$~erg~cm$^{-2}$~s$^{-1}$ exceeds the background, while for the {\\it BeppoSAX} and {\\it ASCA} normalizations it is approximately 80\\% of the XRB. This would rise to 90\\% if we include the extrapolated faint end flux contribution.  The field-to-field differences can clearly be seen in Fig.~\\ref{fig3}, with the CDF-S and SSA22 lying below the CDF-N and SSA13. Between $2\\times 10^{-16}$ and $10^{-14}$~erg~cm$^{-2}$~s$^{-1}$ the CDF-N contributes $0.81\\times 10^{-11}$~erg~cm$^{-2}$~s$^{-1}$~deg$^{-2}$, while the CDF-S contributes $0.54\\times 10^{-11}$. This 40\\% difference is substantially above the Poisson noise expected from the number of sources. \\begin{inlinefigure} \\psfig{figure=fig4.ps,angle=90,width=3.4in} \\vspace{6pt} \\figurenum{3} \\caption{The contribution to the XRB versus flux. The solid boxes are the measured values in the combined sample. The lines show the values from the power-law fits. The open boxes show the CDF-N, the open diamonds the CDF-S, the open upward pointing triangles the SSA13 field, and the open downward pointing triangles the SSA22 field. The individual fields are shown only below $10^{-14}$~erg~cm$^{-2}$~s$^{-1}$ where the error bars are small. } \\label{fig3} \\addtolength{\\baselineskip}{10pt} \\end{inlinefigure}  Given these variations and the uncertainties in the flux calibration of both the counts and the background, we estimate the minimum contribution of the resolved sources to the XRB to be $1.3\\times 10^{-11}$~erg~cm$^{-2}$~s$^{-1}$~deg$^{-2}$. Comparing this number to the maximum measured value of the XRB, $1.8\\times 10^{-11}$~erg~cm$^{-2}$~s$^{-1}$~deg$^{-2}$ (\\markcite{vecchi99}Vecchi et al.\\ 1999), we find a maximum residual of $0.5\\times 10^{-11}$~erg~cm$^{-2}$~s$^{-1}$~deg$^{-2}$. The true value of any diffuse component is likely to be considerably lower."
    },
    "0201/astro-ph0201465_arXiv.txt": {
        "abstract": " ",
        "introduction": "We are closing in on neutron stars both observationally and theoretically. Observationally, a number of masses (M) and a few radii (R) have been measured as well as a number of other properties. Theoretically, modern equation of states (EOS) are more reliable due to precision measurements of nucleon-nucleon interactions, detailed calculations of binding energies of light nuclei and nuclear matter which constrain three-body forces, inclusion of relativistic effects, improved many-body and Monte Carlo methods, etc.  Ultimately, we can exploit the one-to-one correspondence between the EOS and the mass-radius relation of cold stellar object: \\bea P(\\rho)\\quad \\Leftrightarrow \\quad M(R) \\eea Observing a range of neutron star M and R thus reveals the EOS (e.g., pressure P versus density $\\rho$) of dense and cold hadronic matter. Possible phase transitions from nuclear matter to quark matter (either can also undergo superfluid transitions at certain densities and temperatures), hyperon matter, kaon or pion condensates, etc., would also be revealed by an anomalous/kinky function $P(\\rho)$ and $M(R)$. The higher the order of the transition is, the smoother will $M(R)$ be and very accurate observations will thus be necessary. On the other hand, just one accurately measured neutron star mass and radius would already constrain the EOS significantly. Information on the EOS at high baryon densities and the presence or absence of phase transitions could guide us in solving QCD after decades of unsuccessful attempts.   In the following I shall give a brief account of the present status on neutron star observations and theory referring to \\cite{Latt,Balb,annrev,physrep} for longer reviews. Subsequently, I shall attempt to recount the most likely possible phase transition in dense nuclear matter with emphasis on quark matter and its possible color superconducting states, as this is most relevant at this conference. Finally, I shall point to important developments expected in the near future.  \\begin{figure}[htb] \\begin{center} \\epsfig{file=M.eps,height=3.5in} \\caption{Neutron star masses vs.~radius for modern \\cite{apr98} and FPS \\cite{FPS} EOS and strange stars \\cite{Bombaci}. The hatched areas represent the neutron star radii and masses allowed for orbital QPO frequencies 1060~Hz of 4U 1820-30 (vertical lines, \\cite{zss97,miller}) and for burster oscillations of 4U 1636-53 assuming $M\\ge 1.4M_\\odot$ (horizontal lines, \\cite{Nath}) area. Models for glitches in the Vela pulsar constrain masses and radii \\cite{Link} below the full line. The radii of RX J1856-3754 from Refs. \\cite{Pons,Kerk} assumes $M=1.37M_\\odot$.} \\label{MR} \\end{center} \\end{figure} ",
        "conclusions": ""
    },
    "0201/hep-ph0201303_arXiv.txt": {
        "abstract": "Recent data on cosmological variation of the electromagnetic fine structure constant from  distant quasar (QSO) absorption spectra have inspired a more general discussion of possible variation of other constants. We discuss variation of strong scale and quark masses. We   derive the  limits on their relative change from (i) primordial Big-Bang Nucleosynthesis (BBN); (ii) Oklo natural nuclear reactor, (iii) quasar absorption spectra, and (iv) laboratory measurements of hyperfine intervals. ",
        "introduction": "Time variation of  major constants of physics is an old and fascinating topic,  its discussion by many great physicists -- Dirac as the most famous example -- had surfaced many times in the past. Recent attention to this issue was caused by astronomical data which seem to suggest a variation of electromagnetic $\\alpha$  at the $10^{-5}$ level for the time scale 10 bn years, see \\cite{alpha}. The issue discussed in this work is related to it, although indirectly. Instead of looking into atomic spectra and testing a stability of electric charge, we will discuss possible variations of $nuclear$ properties induced by a change in strong and weak scales. We will not go into theoretical discussion of why such changes may occur and how they can be related to modification of electromagnetic $\\alpha$. Our aim is to identify the most stringent phenomenological limitations on such a change,  at (i) a time of the order of few minutes, when the Big Bang Nucleosynthesis (BBN) took place, as well as (ii) at the time of Oklo natural nuclear reactor (1.8 bn years ago), (iii) when quasar radiation has been absorbed in the most distant gas clouds (3-10 bn years ago) and (iv) at the present time.   Mentioning relevant literature we start with the BBN limits on electromagnetic $\\alpha$, obtained in  \\cite{Rubinstein}. The main  results come from variation of late-time nuclear reactions. Because of low temperatures and velocities involved at this stage, those reactions have quite significant suppression due to Coulomb barriers, in spite of the fact that only Z=1-3 is involved. These limits are in the following range \\be \\label{em_limit}   |\\delta \\alpha |^{BBN}/\\alpha <0.02 \\ee  In general, all the models for  time variations of electromagnetic/weak/strong interactions can be divided into two distinct classes, depending on whether it originates in (i) $infrared$ or (ii) $ultraviolet$. The former approach ascribe  variations to some hypothetical interaction of the corresponding gauge bosons with some matter in Universe, such as  vacuum expectation values (VEVs) or  ``condensates'' of some scalar fields. Those typically have zero momentum but can have cosmological time dependence. We would not discuss it: for recent example and references see \\cite{VEVS}.  We would however mention few details from two recent examples of the latter approach, by Calmet and H.~Fritzsch \\cite{Calmet:2001nu} and Langacker, Segre and Strassler \\cite{Langacker:2001td}. Their main assumption   is that Grand Unification \\cite{GQW} of electromagnetic, weak and strong forces holds {\\em at any time}. Therefore, a relation between all three coupling constants exists:   truly modified 2 parameters are in this approach the {\\em unification scale}\\footnote{ One might think that if the GUT scale be used to set units, its variation would be impossible to detect without explicit measurements related to gravity. But it is not so, since the cosmological expansion itself (which is quite important for BBN) contains the Newton's constant (or the Plank mass) in the Hubble constant.  } $\\Lambda_{GUT}$ and {\\em the value of the unified coupling} $\\alpha_{GUT}$ at this scale. Their time variation is assumed to propagate down the scales by the usual (unmodified) renormalization group.  If this assumption is correct, any variation of electromagnetic $\\alpha$ should be accompanied by a  variation of strong and weak couplings as well. Specific predictions need a model, we will mention the one discussed in \\cite{Langacker:2001td}. In it the QCD scale  $\\Lambda_{QCD}$ (determined as usual by a continuation of the running coupling constant into its -- unphysical -- Landau pole) is modified as follows \\be \\label{QCD} {\\delta \\Lambda_{QCD} \\over  \\Lambda_{QCD}}\\approx 34 {\\delta \\alpha \\over \\alpha} \\ee  Another focus of our work is possible limits on cosmological modifications of {\\em quark masses}. According to Standard Model, they are related to electroweak symmetry breaking scale, as well as to some Yukawa couplings $h_i$  . In \\cite{Langacker:2001td}  running of those  has been considered, with a (model-dependent) conclusion that quark mass indeed may have  a different (and stronger) change \\be \\label{mq} {\\delta m_q \\over m_q}\\approx 70 {\\delta \\alpha \\over \\alpha} \\ee Large coefficients in these expressions are generic for GUT and other approaches, in which modifications come from high scales: they appear because weak and strong couplings run more.  If a coefficients of such a magnitude are indeed there, at the BBN time of few minutes the QCD scale and quark masses would be  modified quite a bit, if the upper limit (\\ref{em_limit}) be  used in the r.h.s.  The type of questions we are trying to answer in this work are: Do we know whether it might or might not actually happened? Which simultaneous change of strong and weak interaction scales is or is not  observable? What observables are the most useful ones, for that purpose? What are the actual limits on their variation which can be determined from BBN and other cosmological and laboratory data?  Let us repeat, that although we use the above mentioned papers as a motivation, we do not rely on any particular model. Nevertheless, we will at the end of the paper return to these predictions in order to see whether our limits on time variation imply stronger or weaker effects than the electromagnetic ones. ",
        "conclusions": "Combining our strongest limits on the deuteron binding, from deuteron  (\\ref{lim_d}) and $Li^7$ (\\ref{limLi}), corresponding to variation of their production by a factor 2, with a relation between modification of the deuteron binding and modification of the pion mass (\\ref{modEd}). Both effects suggest about the same BBN limit on the modification of the pion mass {\\em relative to strong interaction scale} $\\Lambda_{QCD}$.  Using eq. (\\ref{modEd}) we obtain: \\be |\\delta_\\pi|_{BBN} < 0.005  \\ee Eq. (\\ref{piEd}) provides a more conservative limit \\be |\\delta_\\pi|_{BBN} < 0.03  \\ee and we think the true limit is somewhere in between.  We have investigated other effects, such as binding of $He^5$ or $pp,nn,np$ (S=0) states, but found that in these cases the needed pion modification about an order of magnitude larger. (It is expected, since all these states  are more loosely bound than the deuteron.) If $m_s$ modification relative to strong scale are as large as our limit on light quark modification just mentioned, it means the nucleon mass can be modified within $\\pm 2 \\, MeV$ due to strange term. Note also, that our limit on quark mass modification is  stronger than the limit \\cite{Langacker:2001td} coming from proton-neutron mass difference (\\ref{npdiff}).  We also pointed out significantly weaker limits on a $simultaneous$ modification of strong scale and $m_q$ scale at the same rate, relative to the gravity scale \\be {\\delta (\\Lambda_{QCD}/M_P) \\over (\\Lambda_{QCD}/M_P)} < 0.1 \\ee Limits on a variation of the quark mass relative to strong scale at the $10^{-4}$ level 3-10 bn years ago follows from observations of distant objects (\\ref{lim_obsrv}), while at the time 1.8 bn years ago the Oklo data lead to even better limits, at the $10^{-8}-10^{-9}$ level.  Although there is no general relation between variation of weak, strong and electromagnetic constants, as we mention in the Introduction it is implied by Grand Unification \\cite{Calmet:2001nu,Langacker:2001td}. If one uses those (\\ref{QCD},\\ref{mq}), one finds that all our limits on relative weak/strong modification are {\\em much more restrictive} than the corresponding limits on the modification of the electromagnetic $\\alpha$. In the case of astronomical observations, in which variation of alpha seems to be seen, one may either soon find the variation of g factors, or rule out  relations between couplings based on Grand Unification idea.  Finally, let us emphasize that our discussion is semi-qualitative in many aspects, and a lot of quantitative work remains to be done. Theory-wise, the most straightforward thing to do is to add modifications directly into the BBN  code, and get more quantitative limits. Although there seem to be no particular problem with the standard BBN at the moment, it is still true that the calculated yields and observations typically differ by 1-2 standard deviations \\cite{Schramm}. Therefore it seem to be worth wile to make a global fit to data with unrestricted modification parameters (like our $\\delta_\\pi$) and see whether zero value would or would not be the best one.  Another challenge to the theory, probably mostly lattice simulations, is to clarify the issue of the dependence of various hadronic parameters on the strange quark mass $m_s$, especially how universal are the derivatives like (\\ref{NssN}) for all hadrons.  Experimental laboratory work and astronomical observations of distant objects can significantly enhance the limits available today, hopefully with a non-zero effect eventually observed.  {\\bf Acknowledgments} One of the authors (ES) is supported by the US Department Of Energy, while the other (VF) is supported by the Australian Research Council. We are grateful to G.E.Brown, V.F. Dmitriev and V.G. Zelevinsky for useful discussions."
    },
    "0201/hep-ph0201245_arXiv.txt": {
        "abstract": "We consider the dynamics of fermions with a spatially varying mass which couple to bosons through a Yukawa interaction term and perform a consistent weak coupling truncation of the relevant kinetic equations. We then use a gradient expansion and derive the CP-violating source in the collision term for fermions which appears at first order in gradients. The collisional sources together with the semiclassical force constitute the CP-violating sources relevant for baryogenesis at the electroweak scale. We discuss also the absence of sources at first order in gradients in the scalar equation, and the limitations of the relaxation time approximation. ",
        "introduction": "The main unsolved problem of electroweak baryogenesis~\\cite{KuzminRubakovShaposhnikov:1985} is a systematic computation of the relevant sources in transport equations. We shall now present a method for controlled derivation of leading CP-violating sources appearing as a consequence of collisions of chiral fermions with scalar particles in presence of a scalar field condensate. We assume the following picture~\\cite{CohenKaplanNelson:1991} of baryogenesis at a first order electroweak phase transition: when the Universe supercools, the bubbles of the Higgs phase nucleate and grow into the sea of the hot phase. For species that couple to the Higgs condensate in a CP-violating manner that CP-violating currents are created at the phase boundary (bubble wall). These currents then bias baryon number violating interactions mostly in the hot (symmetric) phase, where the B-violating processes are unsuppressed. The baryons then diffuse to the Higgs phase, where the B-violating interactions are suppressed, resulting in baryogenesis.  Kimmo Kainulainen~\\cite{Kimmo} has explained how to systematically derive the CP-violating source in the flow term of the kinetic equation for fermions. For details see {\\rm Paper~I}~\\cite{KPSW1}. The source is universal in that its form is independent on interactions. It can be represented as the semiclassical force originally introduced for baryogenesis in two-Higgs doublet models in~\\cite{JoyceProkopecTurok:1994}, and subsequently adapted to the Minimal Supersymmetric Standard Model (MSSM) in~\\cite{ClineJoyceKainulainen:2000}. This problem involves computation of CP-violating sources from charginos, which couple to the Higgs condensate in a manner that involves fermionic mixing. Here we show how to compute the CP-violating source in the collision term that arises at first order in gradients. We work in the simple model of chiral fermions coupled to a complex scalar field {\\it via} the Yukawa interaction with the Lagrangian of the form~\\cite{KPSW1,KPSW2} \\begin{equation} {\\cal L} = i\\bar{\\psi}{\\mathbin{\\partial\\mkern-10.5mu\\big/}} \\psi - \\bar{\\psi}_Lm\\psi_R -  \\bar{\\psi}_Rm^*\\psi_L  + {\\cal L}_{\\rm yu}, \\label{lagrangian0} \\end{equation} where ${\\cal L}_{\\rm yu}$ denotes the Yukawa interaction term \\begin{equation} {\\cal L}_{\\rm yu} = - y\\phi \\bar{\\psi}_L\\psi_R - y \\phi^* \\bar{\\psi}_R\\psi_L, \\label{lagrangian1} \\end{equation} and $m$ is a complex, spatially varying mass term \\begin{equation} m(u) \\equiv y' \\Phi_0 = m_R(u) + i m_I(u) = |m(u)|\\mbox{e}^{i\\theta(u)}. \\label{mass1} \\end{equation} Such a mass term arises naturally from an interaction with a scalar field condensate $\\Phi_0 = \\langle \\hat \\Phi\\rangle$. This situation is realised for example by the Higgs field condensate of a first order electroweak phase transition in supersymmetric models. When $\\phi$ in~(\\ref{lagrangian1}) is the Higgs field the coupling constants $y$ and $y'$ are identical; our considerations are however not limited to this case.   The dynamics of quantum fields can be studied by considering the equations of motion arising from the two-particle irreducible effective action (2PI)~\\cite{CornwallJackiwTomboulis:1974} in the Schwinger-Keldysh closed-time-path formalism~\\cite{Schwinger:1961,ChouSuHaoYu:1985}. This formalism is for example appropriate for studying thermalization in quantum field theory \\cite{BergesCox:2000}.  \\begin{figure}[htb] \\centering \\includegraphics[height=0.6in,width=4.in]{ctp-contour.eps} \\caption{ The Schwinger {\\it closed-time-path} (CTP) used in the derivation of the 2PI effective action~(\\ref{EffectiveAction}). } \\label{fig:ctp} \\end{figure} We are interested in the dynamics of the fermionic and bosonic two-point functions \\beqa S_{\\alpha\\beta}(u,v) &=& -i\\left<T_{{\\cal C}}\\big[\\psi_\\alpha(u)\\bar{\\psi}_\\beta(v)\\big]\\right> \\label{S} \\\\ \\Delta(u,v) &=& -i\\left<T_{{\\cal C}}\\big[\\phi(u){\\phi}^\\dagger(v)\\big]\\right> \\label{Delta} \\eeqa where the time ordering $T_{{\\cal C}}$ is along the Schwinger contour shown in figure~\\ref{fig:ctp}, which is suitable for the dynamics of out-of-equilibrium quantum fields. ",
        "conclusions": "We have studied collisional sources appearing at first order in gradients of a spatially varying mass for the model of chiral fermions interacting with a scalar field {\\it via} a standard Yukawa interaction. This model is relevant for baryogenesis from fermions interacting with the Higgs condensate on growing bubbles at a strongly first order electroweak phase transition. The self-energies have been approximated by one-loop expressions. We have argued that, at first order in gradients, there is no collisional source in the scalar kinetic equation, indicating that baryogenesis sourced by scalar particles is highly suppressed. We have then proven that there is a CP-violating source in the fermionic equation and performed a quantitative analysis of the source. This source, together with the one from the semiclassical force, comprise the relevant sources for baryogenesis at the electroweak scale. Finally, we have argued that the collision term, when modeled in the relaxation time approximation, contains no CP-violating sources. In order to perform a full quantitative assessment of collisional sources, a two-loop analysis of self-energies is required, which is a work in progress."
    },
    "0201/astro-ph0201293_arXiv.txt": {
        "abstract": "We present a {\\it Chandra}/ACIS-I observation of GRO J1744--28. We detected a source at a position of R.A = 17$^h$ 44$^m$ 33.09$^s$ and Dec. = --28\\degr~44$'$ 27.0$''$ (J2000.0; with a 1$\\sigma$ error of $\\sim$0.8 arcseconds), consistent with both {\\it ROSAT} and interplanetary network localizations of GRO J1744--28 when it was in outburst.  This makes it likely that we have detected the quiescent X-ray counterpart of GRO J1744--28. Our {\\it Chandra} position demonstrates that the previously proposed infrared counterpart is not related to GRO J1744--28. The 0.5--10 keV luminosity of the source is $2 - 4 \\times10^{33}$ \\Lunit~(assuming the source is near the Galactic center at a distance of 8 kpc). We discuss our results in the context of the quiescent X-ray emission of pulsating and non-pulsating neutron star X-ray transients. ",
        "introduction": "}  X-ray transients sporadically exhibit very bright outbursts during which their X-ray luminosity can be as high as $10^{36-39}$ \\Lunit. However, most of their time they are in their quiescent state during which they emit X-rays only at a level of $10^{30-34}$ \\Lunit. The mechanisms behind this quiescent X-ray emission are still not understood (see, e.g., Menou et al. 1999; Campana \\& Stella 2000; Bildsten \\& Rutledge 2002).  The most dominant model for the quiescent properties of non-pulsating neutron star transients, which contain a neutron star with a very low magnetic field ($<10^{10}$ Gauss), is the one which assumes that the X-rays below a few keV are due to the cooling of the neutron star after the accretion has stopped (see, e.g., Brown, Bildsten, \\& Rutledge 1998 and references therein).  In the quiescent X-ray spectra of several of those systems an extra power law component above a few keV is present (see, e.g., Asai et al. 1996, 1998; Campana et al. 1998a; Rutledge et al. 2001), but the nature of this component is even more unclear (see, e.g., Campana \\& Stella 2000 for a discussion).  Detailed studies of the quiescent emission from transient X-ray pulsars (with a neutron star magnetic field strength of $>10^{11}$ Gauss) have been inhibited by the lack of detected systems. So far, only two X-ray pulsars have been detected in quiescence, A 0535+26 (Negueruela et al. 2000) and 4U 0155+63 (Campana et al. 2001). Their X-ray luminosities were consistent with the predictions of the cooling neutron star model for the non-pulsating systems, suggesting that also this model might apply to the pulsating ones. However, due to the low statistics of the data a detailed testing of the model could not be performed. To get more insight in the quiescent emission of X-ray pulsars those systems have to be observed with higher sensitivity instruments like {\\it Chandra} or {\\it XMM-Newton}. Furthermore, more quiescent X-ray pulsars have to be detected to determine if all of those systems are consistent with the cooling model or that some systems might have different properties which would suggest alternative X-ray production mechanisms (e.g., accretion down to the magnetospheric radius; Stella et al. 1994; Corbet 1996). A good candidate for such studies is the ``bursting pulsar'' GRO J1744--28.  GRO J1744--28 was discovered in December 1995 with the Burst and Transient Source Experiment (BATSE) aboard the {\\it Compton Gamma-Ray Observatory} ({\\it CGRO}; Fishman et al. 1995; Kouveliotou et al. 1996). The source exhibited rapidly repeating, very bright X-ray bursts which are likely due to accretion disk instabilities (e.g., Lewin et al. 1996). The source also exhibited pulsed emission with a pulsation frequency of 2.1 Hz (Finger et al. 1996). Its bursting and pulsating nature lead to the source being called the bursting pulsar. So far, two major outbursts have been detected, one starting in December 1995 and the other one a year later in December 1996. The latter one ended in April 1997 (see, e.g., Woods et al. 1999).  Augusteijn et al. (1997) likely detected GRO J1744--28 using a {\\it ROSAT} observation performed in March 1996. This observation showed a bright source with a luminosity of $\\sim 2 \\times 10^{37}$ \\Lunit~(for an assumed distance\\footnote{The distance to the source is unknown but the high column density measured by Nishiuchi et al. (1999) and the proximity of the source on the sky to the center of the Galaxy make a large distance likely.} of 8 kpc; 0.1--2.4 keV).  Archival {\\it ROSAT} observations did not detect a source on this position with an upper limit on the luminosity of a few times $10^{33}$ \\Lunit~(0.1--2.4 keV; Augusteijn et al. 1997). This transient nature of the source makes it very likely that indeed the {\\it ROSAT} source is GRO J1744--28 despite that no bursts or pulsations were detected (note that the large {\\it ROSAT} upper limit on the pulsations obtained by Augusteijn et al. 1997 was completely consistent with the strength of the pulsations as measured with BATSE and the {\\it Rossi X-ray Timing Explorer} [{\\it RXTE}]). Localization of GRO J1744--28 by triangulating the data obtained for this source with {\\it Ulysses} and BATSE (part of the interplanetary network [IPN]) resulted in a position (Hurley et al. 2000) which partly overlapped that of the {\\it ROSAT} error circle, confirming the identification of the {\\it ROSAT} source with GRO J1744--28. Cole et al. (1997) and Augusteijn et al. (1997) identified a possible infrared counterpart at the edge of the {\\it ROSAT} error circle, although its enigmatic properties spurred the suggestion that it might be an instrumental artifact (Augusteijn et al. 1997; but see Cole et al. 1997).  Here we present a {\\it Chandra}/ACIS-I observation of the region containing GRO J1744--28. We discovered a weak source near the center of the {\\it ROSAT} error circle, which is likely the quiescent counterpart of the {\\it ROSAT} transient and therefore likely of GRO J1744--28. ",
        "conclusions": "}  We detected a {\\it Chandra} source in the {\\it ROSAT} error circle (Augusteijn et al. 1997) and IPN ellipse (Hurley et al. 2000) of GRO J1744--28, making it likely that we have detected the quiescent X-ray counterpart of GRO J1744--28. The proposed infrared counterpart (Cole et al. 1997; Augusteijn et al. 1997) is not consistent with our {\\it Chandra} position and the nature of this source is unclear (it might be, as suggested by Augusteijn et al. 1997, an artifact, although Cole et al. 1997 could not confirm this).  Our quiescent flux of $2.5 - 5.3 \\times 10^{-13}$ \\funit~(0.5--10 keV; unabsorbed) is consistent with the non-detection of GRO J1744--28 during the archival {\\it ROSAT} observation reported by Augusteijn et al. (1997; using their count rate upper limit for GRO J1744--28, an 0.5--10 keV unabsorbed flux upper limit range can be derived using PIMMS\\footnote{Available at http://heasarc.gsfc.nasa.gov/Tools/w3pimms.html} of $7-10\\times10^{-13}$ \\funit, depending on the spectral model).  One possible mechanism producing the quiescent emission is residual accretion on to the surface of the neutron star. However, Cui (1997) argued (based on the sudden decrease of the pulsation amplitude at times when the 2--60 keV flux of GRO J1744--28, as measured with {\\it RXTE}, was below $\\sim2\\times 10^{-9}$ \\funit) that GRO J1744--28 is in the ``propeller'' regime at relatively low fluxes. In this regime, the magnetic field of the neutron star inhibits accretion on to the neutron star surface. Our {\\it Chandra} flux for GRO J1744--28 is much lower than the critical flux limit given by Cui (1997) strongly suggesting that the source was in the propeller regime during our observation.  Therefore, accretion on to the surface of the neutron star is not likely to be the cause behind the quiescent X-ray emission of GRO J1744--28. Similar conclusions were also reached for the quiescent emission from the transient X-ray pulsars A 0535+26 (Negueruela et al. 2000) and 4U 0155+63 (Campana et al. 2001).  A possible mechanism to produce the observed quiescent X-rays when the source is in the propeller regime might be accretion down to the magnetospheric radius $r_{\\rm m}$ (e.g., Stella et al. 1994; Corbet 1996; Campana et al. 1998b), which is approximately give by $(GM_{\\rm ns})^{-1/7} \\mu^{4/7} \\dot{M}^{-2/7}$, in which $M_{\\rm ns}$ is the neutron star mass, $\\mu$ the neutron star magnetic moment, and $\\dot{M}$ the accretion rate. The luminosity produced will be $L = G M_{\\rm ns} \\dot{M}/r_{\\rm m}$. Using our measured luminosity and the magnetic field strength found by Cui (1997; $\\sim2\\times10^{11}$ Gauss), then $\\dot{M} \\sim 3 - 5 \\times 10^{15}$ g s$^{-1}$ and $r_{\\rm m} \\sim 2 - 3 \\times 10^{8}$ cm.  This radius is larger than the corotation radius $r_{\\rm c} = (GM_{\\rm ns})^{1/3}(P_{\\rm spin}/2\\pi)^{2/3}$ (with $P_{\\rm spin}$ the neutron star spin period), which is $\\sim10^{8}$ cm for GRO J1744--28, and therefore it fulfills the condition that the source should be in the propeller regime. Although a rough estimate\\footnote{Large uncertainties are present in the exact luminosity of the source (due to uncertainties in the spectral shape and in the distance to the source), its magnetic field strength, and in certain unknown constants (like the accretion efficiency), which have been assumed to be roughly 1. A full detailed discussion of those uncertainties is beyond the scope of this {\\it Letter}.}, it indicates that accretion down to the magnetospheric radius might be able to produce the quiescent X-ray luminosity of GRO J1744--28. However, contributions to the X-ray luminosity might also come from several other mechanisms, like accretion onto the neutron star surface (possibly on localized areas such as the magnetic poles) due to the leakage of matter through the magnetospheric barrier.  Another contribution to the quiescent X-ray emission will be the thermal X-ray emission from the neutron star surface which should give a rock bottom lower limit on the X-ray luminosity. Due to the low statistics of our data and the high column density towards GRO J1744--28, we are not able to accurately probe such a thermal component. However, if the temperature of the thermal component is similar to what has been observed for the quiescent spectra of the non-pulsating neutron star transients, then at least the detected emission in our {\\it Chandra} observation is mostly due to an another component, which has its origin probably in one of the above discussed mechanisms.   Progress in our understanding of the quiescent X-ray emission can be made for GRO J1744--28 by obtaining longer {\\it Chandra} or {\\it XMM-Newton} observations of this source. Such observations will be able to constrain its luminosity and spectral shape much better than we are able to due with the present {\\it Chandra} data and will determine if the quiescent spectrum of GRO J1744--28 is similar to that of the quiescent non-pulsating systems, or that a fundamental difference is present. Detecting the pulsations in observations with sufficient time resolution will also constrain the mechanisms for the quiescent X-rays in this system.  Detecting the pulsations in quiescence would also unambiguously proof that the detected source is GRO J1744--28.   {\\it Note added in manuscript:} After submission of our paper, we became aware of the paper by Daigne et al. (2002) who reported on a {\\it XMM-Newton} observation of GRO J1744--28 in quiescence. Their reported position and X-ray flux for the source are consistent with ours."
    },
    "0201/astro-ph0201400_arXiv.txt": {
        "abstract": "We predict the redshift distribution of Gamma-Ray Bursts (GRBs) assuming that they trace the cosmic star formation history. We find that a fraction $\\ga 50\\%$ of all GRBs on the sky originate at a redshift $z\\ga 5$, even though the fraction of the total stellar mass formed by $z\\sim 5$ is only $\\sim 15$\\%.  These two fractions are significantly different because they involve different cosmological factors when integrating the star formation rate over redshift.  Hence, deep observations of transient events, such as GRB afterglows or supernovae, provide an ideal strategy for probing the high-redshift universe.  We caution, however, that existing or planned flux-limited instruments are likely to detect somewhat smaller fractions of high redshift bursts. For example, we estimate that the fraction of all bursts with redshifts $z\\ga 5$ is $\\sim$10\\% in the case of the BATSE instrument, and $\\sim$25\\% in the case of {\\it Swift}. We also show that the intrinsic distribution of GRB durations is bimodal but significantly narrower and shifted towards shorter durations than the observed distribution. ",
        "introduction": "Gamma-Ray bursts (GRBs) are the brightest electromagnetic explosions in the universe (for a recent review, see Piran 2000).  Popular models for their central engine divide into two main classes: (i) the collapse of a massive star to a black hole (BH) (MacFadyen, Woosley, \\& Heger 2001, and references therein); (ii) the coalescence of a binary system involving a neutron star (NS) and a BH or a NS as a companion (e.g.  Eichler et al. 1989; Janka et al. 1999).  The observed association of long-duration GRBs with star forming regions (Djorgovski et al. 2001c, and references therein), and the possible supernova signatures in rapidly-decaying afterglows (Bloom et al. 1999; Kulkarni et al. 2000; Reichart 2001) favors the first class.  Both classes of models associate GRB progenitors with compact objects (BH or NS) that are the end products in the evolution of massive stars. Hence, the GRB formation history is expected to follow the cosmic star formation history (Totani 1997, 1999; Wijers et al. 1998; Blain \\& Natarajan 2000) up to the highest redshifts ($z\\sim 20$) at which the first generation of stars may have formed (Barkana \\& Loeb 2001). GRBs might therefore provide an ideal probe of cosmic star formation at all redshifts that in particular is unaffected by dust obscuration (e.g., Blain \\& Natarajan 2000; Porciani \\& Madau 2001).  In fact, the top-heavy initial mass function (IMF) predicted for the first stars (Bromm, Coppi, \\& Larson 1999, 2002; Abel, Bryan, \\& Norman 2000, 2002, Nakamura \\& Umemura 2001) favors massive stars which are the likely source of GRB progenitors.  GRB afterglows provide a unique probe of the high redshift universe (Lamb \\& Reichart 2000; Ciardi \\& Loeb 2000). The bright, early optical-UV luminosity of a GRB afterglow is expected to outshine its host galaxy, even more so at high redshifts when the typical galaxies are less massive than their present-day counterparts (Barkana \\& Loeb 2001).  The broad-band afterglow spectrum extends into the far UV and so the absorption features imprinted on it by the intervening intergalactic medium (IGM) can be used to infer the evolution of the neutral hydrogen fraction and the metal abundance of the IGM during the epoch of reionization. In difference from galaxies and quasars, which fade rapidly with increasing redshift due to the increase in their luminosity distance, GRB afterglows maintain an almost constant infrared flux with increasing redshift at a fixed time lag after the GRB trigger in the observer frame (Ciardi \\& Loeb 2000). This follows from the cosmological time--stretching of the afterglow transient (which is intrinsically brighter at earlier times) and from a favorable $K$-correction in the afterglow spectrum.  The {\\it Swift} satellite\\footnote{See http://swift.gsfc.nasa.gov/}, planned for launch in 2003, is expected to localize roughly one GRB per day. Sorting out the subset of all GRBs which originate at high redshifts ($z\\ga 5$) would be of particular interest.  Observers may employ a simple strategy for this purpose.  Photometric data from a small telescope should be used at first to identify those GRBs which possess a Ly$\\alpha$ trough at a wavelength of $0.73\\mu{\\rm m}(1+z)/6$ due to absorption by the IGM. Follow-up spectroscopy of those GRBs could then be done on a 10-m class telescope.  In designing this observing strategy it is important to forecast which fraction of all GRBs originate from different redshifts. For example, it would be impractical to search for those very high redshifts which amount to a fraction smaller than $10^{-3}$ of all GRBs, because barely a single one of them would be found by {\\it Swift} over several years of operation.  In this paper, we use existing observational and theoretical work on the cosmic star formation history to predict the fraction of all GRBs that are expected to originate at different redshifts. In order to keep our results general, we make predictions about {\\it all} GRBs without reference to the detection threshold or redshift horizon of any particular instrument.  To ascertain, however, what the BATSE and {\\it Swift} instruments are expected to detect, we in addition estimate the redshift distributions for these flux-limited surveys.   In \\S 2, we calculate the collapsed fraction of baryons as a function of redshift based on the Press-Schechter formalism, and infer the corresponding redshift distribution of GRBs.  In \\S 3 we use the inferred redshift distribution of GRBs to convert the observed distribution of GRB durations into the corresponding intrinsic distribution, under the simple assumption that its normalized form is redshift independent. Finally, we discuss the implications of our results in \\S 4. ",
        "conclusions": "We have derived the redshift distribution of GRBs out to $z\\ga 20$ under the assumption that the GRB event rate traces the cosmic star formation rate. We find that $\\ga 50$\\% of all GRBs on the sky originate from a redshift of 5 or higher. On the other hand, the fraction of baryons that have been incorporated into stars by $z\\sim 5$ is much smaller, comprising only $\\sim 15$\\% of the stellar mass formed by today.  The difference between the two fractions follows from the different cosmological factors in the redshift integrations for the statistics of transient events on the sky as compared to the census of fossil objects in the local universe.  The favorable statistical bias towards high-redshift events on the sky is expected to apply also to Type II supernova explosions which are related to the formation of massive stars in a similar way as GRBs. Despite their dimming with increasing redshift, high--redshift supernovae will be detectable with sufficiently sensitive telescopes such as the {\\it Next Generation Space Telescope}\\footnote{http://ngst.gsfc.nasa.gov/} (NGST; Miralda-Escud\\'e \\& Rees 1997; Woods \\& Loeb 1998).  In fact, our calculation implies that without any additional bias (such as redshift-dependent dust extinction) approximately half of all Type II supernovae detected by NGST will originate at $z\\ga 5$.  Deep observations of high--redshift GRBs and supernovae offer an ideal window into the earliest epoch of cosmic structure formation. The lengthening of the duration of these transients by a factor $(1+z)$ makes it easier for observers to monitor their lightcurves.  Different instruments may find GRBs up to different redshifts, depending on their detection sensitivity and the highly uncertain GRB luminosity function (Schaefer et al. 2001; Schmidt 2001; Norris 2002). A trigger-unbiased way to infer the redshift evolution of the GRB event rate is to compare the number counts of GRBs with the same absolute (intrinsic) luminosity in different redshift bins.  If future observations of this type were to determine a mean redshift for the GRB distribution significantly lower than the one predicted in this paper, then this would indicate either that GRB formation at high $z$ is substantially suppressed, or that GRBs originate from the coalescence of binaries with a time delay of a few Gyr between the formation of a massive star and the GRB event.  Recent observations indicate that a large fraction, $\\sim 50$\\%, of all well-localized GRBs have no associated optical afterglow, and are classified as ``(optically) dark GRBs'' (e.g., Piro et al. 2002). According to our model, a substantial fraction of these dark bursts could originate from $z\\ga 6$.  The intervening, partially neutral IGM would efficiently absorb the rest-frame UV afterglow that would otherwise have been redshifted into the optical band."
    },
    "0201/astro-ph0201366_arXiv.txt": {
        "abstract": "{The reported nightly mean value of the circular polarization of optical emission observed from the close binary system \\a is $0.06\\% \\pm 0.01\\%$. We discuss the possibility that the observed polarized radiation is emitted mainly by the white dwarf or its vicinity. We demonstrate  that this hypothesis is rather unlikely since the contribution of the white dwarf to the optical radiation of the system is too small. This indicates that the polarimetric data on \\a cannot be used for the evaluation of the surface magnetic field strength of the white dwarf in this system. ",
        "introduction": "Measurements of circularly polarized emission from the close binary system \\a have been reported by three independent groups of authors. Cropper (\\cite{c86}) estimated the circular polarization in the optical to be at the level of $0.05\\% \\pm 0.01\\%$. Similar value (0.06\\%) have been derived by Stockman \\e (\\cite{stock92}). Beskrovnaya \\e (\\cite{bibs}) reported a value for the nightly mean of the circular polarization in the $V$ and $V+R$ passbands to be $p_{\\rm mean}=0.06\\% \\pm 0.01\\%$.  The possible origin of the circularly polarized emission from \\a was first discussed by Bastian \\e (\\cite{bdc88}) in the framework of the {\\it oblique rotator} model (Patterson \\cite{p79}). Within this model (which was widely accepted before 1994), \\a is considered as an {\\it intermediate polar} in which the radiation of the white dwarf is powered by the accretion of plasma onto its surface with the rate of $\\sim 10^{16}\\,{\\rm g\\,s^{-1}}$. Assuming the observed polarized emission to be generated in the shock at the base of the accretion column due to the cyclotron mechanism, Bastian \\e (\\cite{bdc88}) estimated the magnetic field strength of the white dwarf to be in excess of $10^6$\\,G.  This result, however, turned out to be one of the first serious arguments against the accretion-powered white dwarf model of AE~Aqr: If the surface field strength of the white dwarf in \\a is indeed $B_* \\ga$\\,1\\,MG, its magnetospheric radius, \\be\\label{rm} \\r \\simeq 7.7\\,10^9\\ \\kappa_{0.5} \\dot{M}_{16}^{-2/7} M_{0.8}^{-1/7} R_{8.8}^{12/7} \\left(\\frac{B_*}{10^6\\,{\\rm G}}\\right)^{4/7} {\\rm cm}, \\ee is larger than the corotation radius, \\be\\label{rcor} R_{\\rm cor} = 1.5\\,10^9 M_{0.8}^{1/3} P_{33}^{2/3}\\,{\\rm cm}, \\ee by more than a factor of five. Here $\\dot{M}_{16}$ is the mass accretion rate expressed in units of $10^{16}\\,{\\rm g\\,s^{-1}}$, $\\kappa_{0.5}=\\kappa/0.5$ is  accounting for the geometry of the accretion flow\\footnote{The value of this parameter lies within the interval $0.5 \\la \\kappa \\la 1$, where $\\kappa = 0.5$ corresponds to the disk geometry of the accretion flow and $\\kappa = 1 $ -- to the spherical geometry (see e.g. Ghosh \\& Lamb \\cite{gl79})}, and $M_{0.8}$, $R_{8.8}$ and $P_{33}$ are the mass, radius, and the spin period of the white dwarf expressed in units of $0.8 M_{\\sun}$, $10^{8.8}$\\,cm and 33\\,s, respectively (for the system parameters see Table~1 in Ikhsanov \\cite{i00}). Under these conditions, the white dwarf is in the centrifugal inhibition regime and a steady accretion process onto its surface does not occur.  The hypothesis that the accretion power is not responsible for the emission of the white dwarf in \\a got serious grounds after 1994. It was recognized that the intensity of radiation emitted from the polar caps of the white dwarf does not correlate with the flaring in the system (Eracleous \\e \\cite{erac94}). Furthermore, the X-ray spectrum of \\a is soft and essentially differs from the hard X-ray spectra of intermediate polars (Clayton \\& Osborne \\cite{co95}). Finally, the discovery of the rapid spin down of the white dwarf (de~Jager \\e \\cite{jmor94}) and the conclusion that no developed Keplerian accretion disk exists in the system (Wynn \\e \\cite{wkh97}; Welsh \\e \\cite{whg98}) have left no doubts that \\a is not compatible with the {\\it oblique rotator} model and that no intensive plasma accretion onto the surface of the white dwarf occurs (for discussion see Ikhsanov \\cite{i01}). Therefore, the basic assumptions used by Bastian \\e (\\cite{bdc88}) for the interpretation of the circularly polarized emission of \\a are unacceptable and the origin of the polarized radiation in this system remains an unresolved problem.  In this paper we are investigating the validity of the suggestion that the observed circularly polarized radiation originates from the white dwarf or its vicinity. ",
        "conclusions": "We have shown that the white dwarf cannot be considered as the main source of the circularly polarized optical emission detected from AE~Aqr. The basic argument is a very small value of the white dwarf contribution to the visual light of the system. As a consequence, the circularly polarized radiation emitted by the white dwarf is significantly diluted and the expected value of the pulse-averaged circular polarization proves to be essentially smaller than the observed value.  Why was this argument  not taken into account in previous investigations\\,? The reason was the widely accepted but wrong assumption that the system radiation is powered by the accretion of material onto the white dwarf surface. Within this assumption, the contribution of the white dwarf to the system radiation proves to be overestimated by almost an order of magnitude and thus, the problem discussed in this paper simply does not arise. The incorrectness of this approach has been recognized only a few years ago, mainly due to the discovery of the rapid spin-down of the white dwarf (de~Jager \\e \\cite{jmor94}), detailed investigations of the optical/UV properties of the 33\\,s coherent oscillations (Eracleous \\e \\cite{erac94}) and the reconstruction of the diskless mass transfer picture (Wynn \\e \\cite{wkh97}).  One of the important consequences of our conclusion is that the observed circularly polarized emission from \\a cannot be used to estimate the surface magnetic field strength of the white dwarf. In the light of modern views on the system, neither the lower limit suggested by Bastian \\e (\\cite{bdc88}) nor the upper limit given by Stockman \\e (\\cite{stock92}) to the magnetic field strength of the white dwarf can be used. In this situation, the value of the white dwarf magnetic field should be estimated using different methods (see e.g. Ikhsanov \\cite{i98}).  Finally, we would like to note that except for the uncertain nightly mean value, no information about the properties of the circularly polarized emission from \\a is currently available. In this situation, the justification of any alternative ideas about the origin of the polarized emission is very complicated and perhaps even impossible, until highly time resolved  polarimetric observations with large signal-to-noise ratio are performed."
    },
    "0201/astro-ph0201199_arXiv.txt": {
        "abstract": "{ He-accreting white dwarfs with sub-Chandrasekhar mass are revisited. The impact of the use of an extended reaction network on the predicted energy production and characteristics of the detonating layers is studied. It is shown that the considered scenario can be the site of an $\\alpha$p-process combined with a p-process and with a variant of the rp-process we refer to as the pn-process. We define the conditions under which the derived distribution of the abundances of the p-nuclides in the ejecta, including the puzzling light Mo and Ru isotopes, mimics the solar-system one. ",
        "introduction": "The outcome of He-accreting sub-Chandrasekhar white dwarfs (WD) has deserved a special attention since the early 80's (e.g. Nomoto 1982, Woosley et al. 1986). Iben \\& Tutukov (1991) have investigated the evolution of a close binary system leading to the formation of a compact CO WD accreting He from a nondegenerate low-mass companion. Limongi \\& Tornambe (1991) concluded that such systems could in some conditions lead to explosive phenomena. Their relatively high estimated frequency, around 0.01 y$^{-1}$ (Iben \\& Tutukov 1991), have drawn attention to their possible connection with the progenitors of Type Ia supernovae.  He-accreting CO WDs are not viewed today as the most likely candidates for such explosions (e.g. H\\\"oflich \\& Khokhlov 1996, Hillebrandt \\& Niemeyer 2000, Branch 2001), but they might well be responsible for some special types of events. In fact, some one-dimensional calculations (e.g. Woosley \\& Weaver 1994 and references therein) have concluded that He-detonations on the considered WDs could well be identified as peculiar supernovae, characterized by rapidly declining light curves with lower maximum luminosities than those reached by C-deflagration Chandrasekhar-mass WDs. These properties are reminiscent of subluminous supernovae like SN 1991bg. Multidimensional simulations have confirmed the onset of the He-detonation, but have revealed significant differences in the central C-ignition which may be triggered by the He-detonation (Livne \\& Arnett 1995, Garcia-Senz et al. 1999).   This Letter limits its focus to some aspects of the surface He detonation which have not deserved much attention up to now. More precisely, we want to test the classical practice of calculating the energy production and associated nucleosynthesis through a nuclear reaction network made of $(\\alpha,\\gamma)$ captures. This approach is obviously unable to treat the production or captures of protons and neutrons in the detonating layers as well as their impact on the energetics and the nucleosynthesis of the He detonation. Concomitantly, we present the first detailed calculation of the synthesis of the nuclides heavier than the iron peak in the considered He detonation. These problems are tackled in the framework of a 1-D model of the He detonation, the details of which are presented in Sect.~2. Section 3 discusses the impact of the use of an extended reaction network on the predicted energy production and characteristics of the detonating layers.  The composition of the ejected material is analyzed in Sect.~4. We demonstrate that the considered scenario can be the site of an $\\alpha$p-process combined with a p-process and with a variant of the rp-process we refer to as the pn-process. We define the conditions under which the derived distribution of the abundances of the p-nuclides mimics the solar-system one. Conclusions are drawn in Sect.~5. ",
        "conclusions": "This Letter presents the first instance of a clear possibility for $\\alpha$p processed material to be ejected into the interstellar medium. In previously proposed sites, like accreting neutron stars associated to  X-ray bursts or accretion disks around black holes (e.g. Schatz et al. 1998), such an ejection is indeed far from being demonstrated. Of course, the global contribution of He-detonating CO-WD to the galactic nuclidic content remains uncertain, but there is reasonable hope for it not to be negligible in view of the predicted frequency of about 0.01 per year for these events. We also find that this galactically `fertile' $\\alpha$p process is accompanied with an efficient p-process and triggers a variant of the rp-process, the pn-process, which develops in the presence of neutrons and with less protons than the classical rp-process. Most of the p-nuclides, including the puzzling light Mo and Ru isotopes, are found to be co-produced in these conditions in relative quantities close to solar. Unfortunately, they are underproduced (except \\chem{78}{Kr}) with respect to the Ca-to-Fe species. The price to pay in order to avoid this difficulty is an increase of the abundances of the seed s-nuclei by a factor of about 100 over their solar values. The astrophysical plausibility of this enhancement remains to be scrutinized in detail, in particular by studying the impact of rotationally induced mixing.  In spite of this problem, we consider that the results presented here are encouraging enough for justifying an extension of our calculations to other situations involving CO-WD of different masses and accretion rates. These additional simulations will be presented elsewhere, along with a detailed discussion of the characteristics and nuclear physics uncertainties of the associated $\\alpha$p, pn- and p-process flows."
    },
    "0201/astro-ph0201150_arXiv.txt": {
        "abstract": "A systematic  study of the  linear thermal stability  of a medium  subject to cooling, self--gravity  and thermal  conduction is carried  out for  the case when  the unperturbed state  is subject  to global  cooling and  expansion. A general,  recursive WKB  solution for  the perturbation  problem  is obtained which can  be applied to a  large variety of  situations in which there  is a separation of  time-scales for  the different physical  processes.  Solutions are explicitly given and discussed for the case when sound propagation and/or self--gravity are the fastest  processes, with cooling, expansion and thermal conduction operating on slower time-scales.  A brief discussion is also added for the solutions in the cases  in which cooling or conduction operate on the fastest time-scale. The general WKB  solution obtained in this paper  permits solving the problem of the effect of thermal conduction and self-gravity on the thermal stability of a globally cooling and expanding  medium. As a result of the analysis, the critical  wavelength (often called  {\\it Field  length}) above  which cooling makes the perturbations unstable against  the action of thermal conduction is generalized to the case of an  unperturbed background with net cooling. As an astrophysical application,  the {\\it generalized Field  length} is calculated for  a  hot ($10^4  -  10^8$  K), optically  thin  medium  (as pertains,  for instance, for  the hot interstellar medium  of SNRs or  superbubbles) using a realistic  cooling  function  and  including  a  weak  magnetic  field.   The stability  domains are compared  with the  predictions made  on the  basis of models for which  the background is in thermal  equilibrium.  The instability domain of the sound waves, in particular, is seen to be much larger in the case with net global cooling. ",
        "introduction": "Thermal instability often appears in astrophysical media which are undergoing a global process of cooling, expansion or contraction.  A non-exhaustive list of examples includes: thermal instability in a cooling flow in a galaxy cluster \\citep{mathews,balbussoker89,balbus91}, in a stellar wind \\citep{balbussoker89}, or in the hot ISM \\citep{begelman}, and structure formation in a collapsing protogalaxy \\citep{fallrees85}. In all those cases, the unperturbed state is not one of thermal or dynamical equilibrium, so that the stability and the time evolution of the perturbations cannot be studied by means of a classical normal--mode Fourier analysis in time. The situation is complicated by other physical processes which can be important for the unperturbed state or for the perturbation itself, like self-gravity, thermal conduction, and magnetic induction and diffusion.  The stability criteria and the evolution of the perturbations directly depend on the relative time-scale of the latter processes with respect to each other and to the time-scales of cooling and expansion of the background in which they are developing.  The astrophysical literature abounds in examples of thermal stability analyses and their application to specific problems, as in \\citet{parker}, \\citet{weymann}, \\citet{field65}, \\citet{gold}, \\citet{defouw}, \\citet{heyvaerts}, \\citet{mathews}, \\citet{flan}, \\citet{balbus86,balbus91}, \\citet{malagoli}, \\citet{balbussoker89}, \\citet{begelman}, \\citet{loew}, \\citet{chun-rosner}, \\citet{burkert}. A milestone was the paper by \\citet{field65}. He considered a non-expanding homogeneous medium in thermal equilibrium (i.e., with cooling exactly compensating heating), including thermal conduction and magnetic fields. The small--perturbation problem is then amenable to a normal-mode Fourier analysis in space and time. Additionally to magnetosonic waves, which can be strongly modified by the thermal processes, he found a non-oscillatory mode (called condensation mode) which may have an isobaric or isochoric character depending on the amount of cooling or conduction possible during one crossing time of a magnetosonic wave across the perturbation. Conduction becomes important at small spatial scales and, in general, has a strongly stabilizing effect. Another important step forward was taken by \\citet{mathews} and \\citet{balbus86}, who considered a time--dependent unperturbed medium.  By using the entropy equation, they obtained stability criteria for the condensation mode in a medium with global cooling, not taking into account conduction, self-gravity or any background stratification. \\citet{balbus86}, in particular, carried out a WKB analysis in space for radial perturbations in an expanding or contracting medium with spherical symmetry. \\citet{balbussoker89} relaxed the condition of homogeneous background and obtained solutions for a stratified medium. \\citet{balbus91} also included a background magnetic field.  The topic of thermal instability in astrophysical media seems to have been gathering momentum in the past several years. In spite of the foregoing, there is to date no {\\it systematic} study of the stability of a magnetized medium subject to cooling, self-gravity and thermal conduction when the unperturbed state itself is undergoing expansion and cooling. The purpose of the present paper is to carry out such a study using, wherever possible, a WKB expansion {\\it in time}. The treatment introduced here permits recovering many results of the previous literature in a unified manner and allows obtaining new ones. The non-magnetic and non-stratified case is considered in the present paper, which deals with the problem of a uniform (but time-dependent) background. However, the results of this paper can also be applied locally to a weakly non-homogeneous background (see sect. \\ref{swkb}). Several characteristic time-scales have to be considered, to wit, those associated with the background expansion, self-gravity, cooling, sound-wave propagation and conduction.  When background cooling and expansion are {\\it not} the fastest processes in the system, a WKB analysis in time can be carried out. General WKB solutions up to arbitrary order are then obtained and applied to the particular cases when the fastest time-scales are those associated with sound wave propagation, with sound and self-gravity simultaneously, or with conduction.  New results obtained in the present paper are as follows.  First, we obtain an expression for the critical length separating stable from unstable domains which generalizes the classical {\\it Field length} to a system which, in the unperturbed state, is cooling (or otherwise) and expanding or contracting globally. \\citet{field65} showed the opposite effects of conduction and cooling on the stability of a system in static and thermal equilibrium in the case when the cooling function, $\\Ll$, fulfills $[\\partial \\Ll/\\partial T]_p < 0 $, and found that below a critical wavelength (the classical {\\it Field length}) conduction stabilizes the system. Our study permits relaxing the condition of thermal equilibrium in the unperturbed state and including in the latter a global expansion or contraction. We may note that, for lack of a more accurate expression, different authors in the past have tentatively used the value calculated by Field even though their problem involved a medium undergoing global cooling \\citep[e.g.,][]{david89,pistinner}. In section \\ref{s_astr}, we calculate the generalized Field length for a medium undergoing a realistic net cooling, and compare it with the classical Field length of a medium in thermal equilibrium with the same cooling but with a balancing constant heating. The results of the present paper also generalize those of \\citet{balbus86}, who considered a medium undergoing global cooling in the unperturbed state, but ignored thermal conduction.  Second, there is in the literature no simultaneous study of the thermal and gravitational linear stability of a medium undergoing net cooling. The combination of both phenomena is important, for example, for the structure formation in a protogalaxy \\citep{fallrees85}, where condensations previously formed by thermal instability in a cooling medium become gravitationally unstable. In section \\ref{s_sougr} we calculate how the acoustic-gravity modes of \\citet{jeans} are modified through conduction and cooling when the background itself is cooling globally. We also show the existence of a hitherto unknown condensation mode in this case. In all these solutions the perturbations can go through different stability regimes: in a cooling medium the generalized Field length and the Jeans length change with time whereas the wavelength of the perturbation remains constant.  There is an alternative to a WKB expansion when trying to solve the problem tackled in this paper. \\citet{burkert} have used a Taylor expansion in time for the reduced problem when there is no thermal conduction, self-gravity or background expansion.  Yet, a Taylor expansion has the large disadvantage of being valid only at times small compared with the characteristic times of evolution of the background. This is in contrast to the WKB approach taken in the present paper, which is valid for arbitrary times as long as the perturbation does not grow beyond the linear regime.  The present paper is organized as follows: in section \\ref{s_pose_problem}, the equations of the problem are presented and the relevant time-scales introduced both for the background and the perturbations.  Section \\ref{swkb} contains a discussion of the application of the WKB method to the present problem and provides a general WKB solution to the linear perturbation equations. In section \\ref{sec_intermediate}, a collection of solutions for intermediate wavelengths is presented. The word {\\it intermediate} is used in the sense that the wavelength is not small nor large enough for conduction or cooling, respectively, to be the dominant processes in the evolution.  The case in which sound propagation is the fastest process is considered in section \\ref{ssou} whereas in section \\ref{s_sougr} both self-gravity and sound are the fastest processes. Section \\ref{sother} is devoted to the large and small wavelength regimes and, in particular, to the cases when cooling or conduction operate on the fastest time-scale. In section \\ref{s_astr}, the generalized Field length and the instability domain are computed for an optically thin medium at high temperature threaded by a magnetic field without dynamical effects (e.g., the hot ISM). The final section (sect.~\\ref{ssum}) contains a summary and further discussion of the results of the paper. ",
        "conclusions": "\\label{ssum}  We have carried out a systematic analysis of the thermal stability of a uniform medium which, in the unperturbed state, is undergoing net cooling and expansion, so that the problem is {\\bf not} amenable to a full Fourier normal mode analysis. Thermal conduction and self-gravity have also been included as fundamental ingredients. The small-perturbation problem yields a system of ordinary differential equations with three independent solutions. In many cases, two of them have the character of oscillatory modes (typically sound waves, or, more generally, acoustic-gravity modes) and the third one is a non-oscillatory (or condensation) solution. The wavelength of the perturbation determines the positions of the different physical processes in the hierarchy of timescales.  At small enough wavelength, thermal conduction damps the condensation mode and makes sound waves isothermal. At large wavelength, one or a few of the following processes are dominant: cooling, self-gravity and expansion.  For example, if cooling is the only dominant process, the perturbation evolving on the fastest timescale is isochoric with no sound waves present that could even out the pressure gradients. At intermediate wavelength, the propagation of sound is the fastest process, so the evolution of the condensation mode is isobaric.  In the paper, solutions have been obtained in a number of instances using the WKB method. WKB solutions can  be found only if two conditions are fulfilled: (1) there is a clear separation of timescales between {\\it fast} and {\\it slow} physical processes in the system, and (2) the background expansion and cooling belong to the {\\it slow} category.   Of special interest in this paper are those solutions whose growth rate is of the same order as the rate of change of the unperturbed background. In this case, the change in time of the critical wavelengths (the generalized Field length, $\\lambcrit$, and the Jeans length, $L_J$) may make the perturbation go through totally different stability regimes during their evolution. For example, the solutions obtained in section \\ref{s_sougr} for a medium undergoing net cooling include a condensation mode initially in the stable range whose wavelength becomes larger than the Field length after some time, thus becoming thermally unstable. Later on, the Jeans length decreases sufficiently so that the perturbation also becomes gravitationally unstable. A sound wave may undergo a similar process of change of stability, first being in a stable regime with slowly decreasing amplitude, then becoming thermally unstable and finally, when the self-gravity frequency becomes comparable with the sound frequency, turning into an acoustic-gravity mode which ends up as a gravitationally-dominated perturbation.   The analysis of the present paper shows (section \\ref{s_astr}) that in a medium undergoing a realistic net cooling the sound waves are thermally unstable for a comparable range of wavelengths and temperatures as the condensation mode. This may come as a surprise: in contrast, and by way of example, if the thermal equilibrium is reached by balancing the cooling with a constant heating, the sound waves are stable for almost all temperatures in the range $10^4$ -- $10^8$ whereas the condensation mode continues being unstable. On the other hand, the inclusion of a weak magnetic field has a strong influence on the thermal conduction. For astrophysical environments such as the hot ISM and the hot IGM, the generalized Field length along the magnetic field is many orders of magnitude larger than in the perpendicular direction. Therefore, the unstable condensations will be filaments directed along the magnetic field lines.  For the description of the astrophysical systems in which thermal instabilities appear, it is important to study the non-linear phase of both condensation mode and sound waves. The unstable sound waves develop pairs of shock fronts and rarefaction waves, as known from elementary hydrodynamics. The condensation mode, instead, just continues growing without propagating. The detailed physics of the nonlinear phase has to be calculated, in general, using numerical means. Yet, a number of results can be advanced at this stage. For instance, in the linear analysis, a perturbation with negative $\\rho_1/\\rho_0$ grows (in absolute value) as fast as a perturbation with positive $\\rho_1/\\rho_0$. This symmetry is broken in the non-linear evolution. In a medium undergoing net cooling, for example, a condensation with positive $\\rho_1/\\rho_0$ (negative $T_1/T_0$) grows faster than one with $\\rho_1/\\rho_0 < 0$. This happens because of the sum of the following two processes. First, thermal conduction grows with temperature. Therefore, a condensation with positive $T_1/T_0$ is more damped by thermal conduction than a negative one. Second, in the range $10^5 - 10^7$ cooling increases toward lower temperatures (Fig.~\\ref{fig_apl1}). Therefore, a condensation with negative $T_1/T_0$ is more destabilized by cooling than a positive one. A related aspect is that the characteristic spatial size (width at half height) of a condensation decreases by orders of magnitude during its evolution. This happens because the coldest zones of the perturbation cool much faster than the warmer ones. The last aspect can be seen in the paper of \\citet{david88}, where the non-linear evolution of a perturbation of this type is computed numerically.  In the weakly--magnetized case, at the beginning of the non-linear evolution ($\\beta \\gg 1$), the compression in the filament proceeds mainly in the directions perpendicular to the magnetic field lines. This is because of the reduction of thermal conduction transversely to the field lines: a much higher temperature (and thus pressure) gradient can be maintained across than along the field lines. The plasma $\\beta$ inside the condensation decreases with time because of the temperature decrease and the growth of magnetic field due to the compression. In the advanced non-linear phase, the perpendicular compression comes to a halt once $\\beta \\ll 1$, since then the net force across the field lines (the gradient of the sum of magnetic and thermal pressures) is almost decoupled from temperature. After this, the compression is only possible along the field lines. As a related example, see \\citet{david89}, which numerically compute the non-linear evolution of a planar condensation taking into account magnetic pressure.  The WKB analysis that we have carried out in this paper assumes a uniform background. However, our results are also applicable for a weakly non-uniform background. More in detail, our results for the condensation mode in the sound domain are valid if the Brunt-V\\\"ais\\\"al\\\"a frequency associated with the background entropy gradient is much smaller than the growth rate of the condensation mode. This holds, for example, in the supersonic region of a wind and inside old SNRs and superbubbles. In contrast, for the study of the thermal instability in cooling flows in clusters of galaxies, the buoyancy of the perturbations must be taken into account. Many authors have studied this problem: \\citet{mathews}, \\citet{malagoli}, \\citet{balbussoker89}, \\citet{loew} and \\citet{balbus91}.  \\citet{loew} and \\citet{balbus91}, in particular, conclude that a magnetic field, even a weak one, can inhibit buoyancy and trigger thermal instability, which otherwise would be small due to buoyancy. The extension of the results of the present paper to environments of this type requires additional work."
    },
    "0201/astro-ph0201220_arXiv.txt": {
        "abstract": "We present a rapid binary evolution algorithm that enables modelling of even the most complex binary systems. In addition to all aspects of single star evolution, features such as mass transfer, mass accretion, common-envelope evolution, collisions, supernova kicks and angular momentum loss mechanisms are included. In particular, circularization and synchronization of the orbit by tidal interactions are calculated for convective, radiative and degenerate damping mechanisms. We use this algorithm to study the formation and evolution of various binary systems. We also investigate the effect that tidal friction has on the outcome of binary evolution. Using the rapid binary code, we generate a series of large binary populations and evaluate the formation rate of interesting individual species and events. By comparing the results for populations with and without tidal friction we quantify the hitherto ignored systematic effect of tides and show that modelling of tidal evolution in binary systems is necessary in order to draw accurate conclusions from population synthesis work. Tidal synchronism is important but because orbits generally circularize before Roche-lobe overflow the outcome of the interactions of systems with the same semi-latus rectum is almost independent of eccentricity. It is not necessary to include a distribution of eccentricities in population synthesis of interacting binaries, however, the initial separations should be distributed according to the observed distribution of semi-latera recta rather than periods or semi-major axes. ",
        "introduction": "\\label{s:intro}  The evolution of binary stars does not differ from that of single stars unless they get in each other's way. If the binary orbit is wide enough the individual stars are not affected by the presence of a companion so that standard stellar evolution theory is all that is required to describe their evolution. However if the stars become close they can interact with consequences for the evolution and appearance of the stars as well as the nature of the orbit.  The effective gravitational potential in a frame rotating with a circular binary system forms equipotential surfaces called Roche surfaces. A sphere of the volume enclosed by the critical Roche surface defines the Roche-lobe radius of each star. If either star fills its Roche-lobe then gas flows from the outer layers of that star through the inner Lagrangian point that connects the two Roche-lobes. Some or all of this gas may be captured by the companion star so that mass transfer occurs and, as a result, the subsequent evolution of both stars takes a different course from that of isolated stars. When the Roche-lobe filling star is a giant, with a convective envelope, or is significantly more massive than its companion then, as described by Paczy\\'{n}ski (1976), the transferred mass may not be captured by the companion but instead accumulates in a common envelope surrounding both stars. The outcome of common-envelope evolution is still not fully understood but possible scenarios include loss of the envelope as the two cores spiral-in to form a closer binary or coalescence of the two stars.  Even in a detached system it is still possible for the stars to interact tidally. Tides can synchronize the spin of the stars with the orbit and circularize an eccentric orbit as the binary tends towards an equilibrium state of minimum energy. Further, if one star is losing mass in a stellar wind the companion may accrete some of the material with consequences for the orbit. For a discussion and review of the processes involved in close binary evolution and the various kinds of binaries or exotic stars that can result see Pringle \\& Wade (1985) and Wijers, Davies \\& Tout~(1996).  The effects of close binary evolution are observed in many systems, such as cataclysmic variables, X-ray binaries and Algols, and in the presence of stars such as blue stragglers which cannot be explained by single star evolution. While many of the processes involved are not understood in detail we do have a qualitative picture of how binaries evolve and can hope to construct a model that correctly follows them through the various phases of evolution. Initial conditions are the mass and composition of the stars, the period (or separation) and eccentricity of the orbit. In order to conduct statistical studies of complete binary populations, i.e. population synthesis, such a model must be able to produce any type of binary that is observed in enough detail but at the same time be computationally efficient. By comparing results from the model with observed populations we can enhance our understanding of both binary evolution and the initial distributions (e.g. Eggleton, Fitchett \\& Tout 1989; Tutukov \\& Yungelson 1996; Terman, Taam \\& Savage 1998; Nelson \\& Eggleton 2001).  Models for binary evolution have been presented in the past (e.g. Whyte \\& Eggleton 1985; Pols \\& Marinus 1994; Portegies Zwart \\& Verbunt 1996). The model we present here supersedes the work of Tout et al. (1997) primarily by including eccentric orbits and stellar spins, which are subject to tidal circularization and synchronization. Amongst other improvements the possibility of mass accretion from a wind is included. Our model incorporates the detailed single star evolution (SSE) formulae of Hurley, Pols \\& Tout (2000, hereinafter PapI) which allow for a wider range of stellar types than the description of stellar evolution used by Tout et al.\\ (1997). This requires an updating of the treatment of processes such as Roche-lobe overflow, common-envelope evolution and coalescence by collision.  Throughout this paper we refer to one star as the primary, mass $M_1$, stellar type $k_1$ etc., and the other as the secondary (or companion), mass $M_2$ and stellar type $k_2$. At any time the primary is the star filling, or closest to filling, its Roche-lobe. Numerical values of mass, luminosity and radius are in solar units unless indicated otherwise. The stellar types correspond to the evolutionary phases desiginated by the rapid SSE algorithm of PapI, which are:  \\begin{eqnarray*} 0 & = & \\mbox{ MS star } M \\la 0.7 \\mbox{ deeply or fully convective} \\\\ 1 & = & \\mbox{ MS star } M \\ga 0.7 \\\\ 2 & = & \\mbox{ Hertzsprung Gap (HG)} \\\\ 3 & = & \\mbox{ First Giant Branch (GB)} \\\\ 4 & = & \\mbox{ Core Helium Burning (CHeB)} \\\\ 5 & = & \\mbox{ Early Asymptotic Giant Branch (EAGB)} \\\\ 6 & = & \\mbox{ Thermally Pulsing AGB (TPAGB)} \\\\ 7 & = & \\mbox{ Naked Helium Star MS (HeMS)} \\\\ 8 & = & \\mbox{ Naked Helium Star Hertzsprung Gap (HeHG)} \\\\ 9 & = & \\mbox{ Naked Helium Star Giant Branch (HeGB)} \\\\ 10 & = & \\mbox{ Helium White Dwarf (HeWD)} \\\\ 11 & = & \\mbox{ Carbon/Oxygen White Dwarf (COWD)} \\\\ 12 & = & \\mbox{ Oxygen/Neon White Dwarf (ONeWD)} \\\\ 13 & = & \\mbox{ Neutron Star (NS)} \\\\ 14 & = & \\mbox{ Black Hole (BH)} \\\\ 15 & = & \\mbox{ massless remnant.} \\end{eqnarray*}  The SSE algorithm provides the stellar luminosity $L$, radius $R$, core mass $M_{\\rm c}$, core radius $R_{\\rm c}$, and spin frequency $\\Omega_{\\rm spin}$, for each of the component stars as they evolve. A prescription for mass loss from stellar winds is included in the SSE algorithm. The algorithm covers all the evolution phases from the zero-age main-sequence (ZAMS), up to and including the remnant stages, and is valid for all masses in the range 0.1 to $100 \\Msun$ and metallicities from $Z = 10^{-4}$ to $Z = 0.03$. This rapid binary evolution algorithm is a natural extension of the SSE algorithm. Many of the formulae and much of the terminology contained in PapI are utilised in this current paper and therefore the interested reader is encouraged to review PapI.  In Section~2 we describe the binary evolution algorithm in detail. Section~3 contains illustrative examples of Algol and cataclysmic variable evolution and also compares our model with the results of certain binary cases highlighted by Tout et al. (1997) and other authors. We present the results of population synthesis to examine the effects of tides on binary evolution in Section~4 and conclusions are given in Section~5. ",
        "conclusions": "\\label{s:bconc}  The primary purpose of this paper was to describe in detail a rapid evolution algorithm for binary stars. Recognizing that tidal interaction is an important process in binary evolution we have also studied its effect on the results of binary population synthesis and made a global comparison with observations.  The main effects of tides on the evolution of an interacting binary are: \\begin{itemize} \\item[(a)] circularization of the orbit, generally well before RLOF occurs, and \\item[(b)] the exchange of angular momentum between the orbit and the spins of the components. \\end{itemize} Because as a star expands tidal synchronization transfers orbital angular momentum to the stellar rotation, in general the latter effect tends to bring binary components closer together and cause them to follow an evolutionary path similar to that of a closer binary if tides were ignored. However, we find that as a result of the sometimes very convoluted evolutionary paths of interacting binaries, this general principle does not always hold and it is difficult to summarize the overall effect of tides on binary evolution.  The primary differences in our population synthesis results when tidal evolution within binary systems is ignored are as follows: \\begin{enumerate} \\item a 35\\% decrease in the birth rate of classical CVs but a 23\\% increase in sdB CV production; \\item lower birth rates for LMXBs with the persistent NS LMXB rate decreasing significantly but conversely the number of NS LMXB increases; \\item a 10\\% increase in the birth rate of double-degenerates; \\item an increase of 190\\% in the incidence of exploding HeWDs; \\item and 55\\% more Algols currently in the Galaxy. \\end{enumerate} While these differences may seem substantial they are generally not enough to confirm one way or another whether tidal evolution helps to explain the observed binary populations. This is mainly due to the lack of comprehensive binary searches and the selection effects involved that make discrimination on the basis of observational evidence a risky business. However, tides are present in binary systems whether we like it or not. One area where neglecting tidal evolution does seem in clear conflict with observations is in the incidence of HeWDs that explode as supernovae. This rate is too high by a factor of two to be accounted for by either the observed numbers of type Ia or type Ib SNe (see Tout et al.~2001 for further discussion).  The common-envelope efficiency parameter ${\\alpha}_{\\rm\\SSS CE}$ is an uncertain factor in a phase of evolution that is crucial for the production of many types of binary. The failure of the model with ${\\alpha}_{\\rm\\SSS CE} = 1$ to produce anywhere near enough DDs would seem to favour ${\\alpha}_{\\rm\\SSS CE} = 3$. Choosing the secondary mass independently of the primary mass does not produce enough DDs, symbiotics or type Ia SN candidates. All in all the results favour using the properties of Model~A. In fact the standard tidal model, represented by Model~A, is not in disagreement with any of the observations, except in the production of too many type Ib/c SNe. This discrepancy could indicate a lower mass loss rate for helium stars.  A major conclusion of this work must be that it is extremely difficult to set contraints on any of the parameters involved in binary evolution from population synthesis of birth rates and galactic numbers of the various types of binary. This is certainly true while such a large number of parameters remain uncertain. The task would be simpler if we could find observational tests specific to an individual parameter. For example, Algols do not require common-envelope evolution for formation. Therefore, if observational constraints on the number of Algol systems currently in the Galaxy improve then this could provide a suitable test for models of thermal-timescale mass transfer. The effect of varying additional model parameters, such as the binary enhanced mass loss, and initial conditions, such as the separation distribution, would need to be quantified before such a test could be reliably implemented. This is beyond the scope of this paper but will be the subject of future work. In terms of tides it is difficult to isolate direct tests of the tidal evolution model with the type of population synthesis performed in this paper. In fact, to properly constrain the strength of tidal interaction it is necessary to study pre-RLOF binaries, for example, as has been done previously by Zahn (1975, 1977). What we would like to emphasise is that the outcome of binary evolution is sensitive to the physical processes of tidal circularization and synchronization. Therefore, any attempt to constrain uncertain parameters in binary evolution by the method of population synthesis must utilise a binary evolution algorithm that incorporates a working model of tidal evolution, such as that presented here.  An exciting challenge for the future involves attempting to reproduce the individual orbital characteristics of a large number of observed binaries with our binary evolution model. This will allow multi-parameter fitting and should become a powerful tool as the statistics of binary surveys continue to improve. Our preliminary work in this area has shown that the observed properties of DD binaries are much easier to explain when tidal evolution is included. A detailed exploration of the DD parameter space will be the focus of another paper.  Tidal synchronism is important but because orbits generally circularize before Roche-lobe overflow the outcome of the interactions of systems with the same semi-latus rectum is almost independent of eccentricity. Although the inclusion of a distribution of eccentricities seems natural it is not necessary in population synthesis of interacting binaries, however, the initial separations should be distributed according to the observed distribution of semi-latera recta rather than periods or semi-major axes.  A necessity for the near future is a thorough comparison of our BSE algorithm with the workings of a detailed evolution code (e.g. Nelson \\& Eggleton 2001). In this way we can improve the algorithm, especially the treatment of the Hertzsprung gap phase, and hopefully add more stringent constraints to many of the evolution variables. Work is already underway to provide more accurate descriptions of the parameters $k'_2$ (see Section~\\ref{s:beqtid}) and $\\lambda$ (see Section~\\ref{s:comenv}), through investigation of the detailed stellar models provided by Pols et al.~(1998). The algorithm will also be improved by providing options for how angular momentum is lost from the binary system during non-conservative mass-transfer (see Section~\\ref{s:rlof}), reflecting the various possible modes of mass-transfer (Soberman, Phinney \\& van den Heuvel 1997). When this work is completed we will perform additional population synthesis calculations in order to quantify how the various improvements affect the results presented here."
    },
    "0201/gr-qc0201051_arXiv.txt": {
        "abstract": "We show that a reformulation of the ADM equations in general relativity, which has dramatically improved the stability properties of numerical implementations, has a direct analogue in classical electrodynamics.  We numerically integrate both the original and the revised versions of Maxwell's equations, and show that their distinct numerical behavior reflects the properties found in linearized general relativity.  Our results shed further light on the stability properties of general relativity, illustrate them in a very transparent context, and may provide a useful framework for further improvement of numerical schemes. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0201/astro-ph0201528_arXiv.txt": {
        "abstract": "Various implications of new, non-perturbative pomeron inspired enhancement of small-$x$ \\nN\\ \\sfts\\ for high-energy \\n\\ astrophysics are discussed. At $x \\gtrsim 10^{-5}$ these functions are given by perturbative \\emph{QCD}, while at lower $x$ they are determined by a specific generalization of $F_2^{ep}(x,Q^2)$ description, proposed by A. Donnachie and P. V. Landshoff (their two-component model comprises \\emph{hard} and \\emph{soft} pomerons), to \\nuN-scattering case.  We found that i) such enhancement causes the most rapid growth of $\\nu N$-\\css\\ at high energies, ii) that pomeron effects may be perceptible in the rates of \\n\\ induced events in future giant detectors and iii) that the rate of high-energy \\n\\ flux evolution (due to absorption ($CC+NC$) and regeneration ($NC$)) on its pass through a large column depth of matter may be subjected to additional influence of hard pomeron. Solving transport equations for the initially power-law decreasing $\\nu$-spectra, we have evaluated shadow factors for several column depths and spectrum indices. The results are compared with analogous calculations, performed within a trivial small-$x$ extrapolation of \\sfts. Hard pomeron enhanced high-energy shadow factors are found to be many orders of magnitude lower than those obtained within ordinary perturbative \\QCD. ",
        "introduction": "\\label{introduction}% As a rule, accelerator experiments provide data of very high quality. But such measurements seem to be hardly possible at energies, $E \\gtrsim 1 \\times 10^{19}$~eV, even in the very distant future.  On the other hand, \\emph{non-accelerator}, observational physics gives a good chance to get data at ultrahigh energies \\cite{BBGGP}. But the accuracy of data in this case is worse  due to scarcer statistics and lower resolution of detectors. It is remarkable, that non-accelerator \\emph{High-Energy (HE)} physics combines into one the 'telescope' and 'microscope' physics. Really, the detected \\emph{UltraHigh-Energy (UHE)} particles are produced in perhaps most distant, extragalactic cosmic sources. Hence, it deals with extremely large, cosmological, distances. On the other hand, interactions of such particles with matter occur at least distances.  Realization of \\UHE\\ astrophysics is a complex problem. It poses many hard questions, such as:  \\vspace{1mm}% 1. what types of ultrahigh-energy particles may be observed?  \\vspace{1mm}% 2. where and how ultrahigh-energy particles are produced?  \\vspace{1mm}% 3. what detectors are to be used?  \\vspace{1mm}% 4. what is the expected rate of events in a detector?  \\vspace{1mm}% 5. does this rate exceed the background?  \\medskip Let us try to answer these questions in brief. \\subsection{Particles} Through the centuries people used only star light, viz.\\ optic photons, for the sky observation. Today actually the whole range of e.m.\\ radiation, up to TeV photons \\cite{KASKADE}, is involved in measurements.  Another well-known example of high-energy cosmic radiation is provided by Cosmic Rays \\emph{(CRs)}. These are mostly protons ($H$), $He$, $C$, $N$ and up to $Fe$ ions, which continuously bombard the atmosphere. Energy spectrum of \\CRs\\ extends from $1$~GeV up to perhaps $1 \\times 10^{21}$~eV; the highest energy event with $E \\simeq 3\\times 10^{20}$~eV has been detected by the \\emph{AGASA} array \\cite{AGASA}. The nature of incident particles in such events as well as where and how they are produced is unknown. Moreover, the average path length of $E \\gtrsim 6 \\times 10^{19}$~eV protons in the cosmic microwave background radiation (with temperature $T=2.73K$) is essentially limited by $ p + \\gamma \\rightarrow \\pi^+ + n$ interactions \\cite{GZK} (\\emph{GZK} cut-off), so that their presence in \\CRs\\ is a problem.  Besides traditional photons and \\CRs, gravitation waves and \\ns\\ are also regarded as possible instruments for \\HE\\ astrophysics. In this paper we concentrate on \\ns, in particular, on \\UHE\\ \\ns\\ with $E_\\nu \\gtrsim 1\\times 10^{-15}$~eV. Being neutral, they do not deflect in magnetic fields, hence their arrival directions shoot back to the production site. Small \\css\\ allow these \\ns\\ to travel extremely large distances without absorption. But their production in cosmic sources is difficult (i.e.\\ sources are rare) and these sources are distant. It makes the expected flux to be small. Moreover, detection of such \\ns\\ is complicated by the same small \\css. To compensate for these negative factors, one should construct gigantic detectors. \\subsection{Production} There are many theoretical models predicting large luminosity of cosmic sources in high-energy particles, including \\ns\\ (see e.g.\\ \\cite{BBGGP}). For example, such fluxes may be produced in supernova explosions, in active galactic nuclei, during gamma-ray bursts etc. Energy outburst in these sources, including in the form of \\HE\\ particles, may be really high, but the expected flux at the Earth remains low due to large distance.  Basically, there are two models of \\HE\\ particle production. First assumes proton acceleration up to very high energies on cosmic source shocks (usually gas accretion onto a massive black hole, presumably residing in center of an active galaxy, is suggested), sometimes in relativistic jets pointing to the Earth. Then pions and kaons appear as a result of $pp$- and/or $p\\gamma$-collisions in a source. These mesons decay, giving birth to \\HE\\ \\ns. The described scenario is called \\emph{down-top} mechanism. Unfortunately, it is difficult to obtain proton energies higher than $E \\sim 10^{17}$~eV in this mechanism; one needs both large shock radius and high magnetic field.  Another possibility is proposed by the \\emph{top-down} mechanism (for review and discussion of upper limits to the expected $\\nu$-fluxes see e.g.\\ Ref.~\\cite{NPQSGAZ} and references therein). It is based on suggestion that some yet unknown, extremely massive (with $m_X \\sim 10^{14} \\div 10^{16}$~GeV) long-living gauge $X$-bosons decay to the ordinary particles. These bosons allegedly originate from topological defects, which in their turn have been produced during a hypothetical phase transition in the early Universe. A big fraction of $X$-boson's energy goes to \\HE\\ particles (\\ns). These attractive models involve new physics beyond Standard Model and seem a bit speculative. An experimental verification is needed. \\subsection{Detectors} \\HE\\ \\n\\ detectors should be very large and well shielded from background radiations. The are to have either $S\\gtrsim 1$~km$^2$, especially if secondary muons (say from $\\nu_\\mu + N \\rightarrow \\mu^- + X$) are detected, and/or gigantic mass, $M \\gtrsim 1\\times 10^9$~ton, in the case of registration using nuclear-electromagnetic cascades (from hadron state $X$ in the $\\nu_\\mu$ case or from $\\nu_e + N \\rightarrow X$). These requirements may be met: \\begin{itemize} \\item by search for Cerenkov radiation in deep underwater detectors with desirable volume under control $V \\sim 1$~km$^3$ (mass $M \\sim 10^9$~ton); \\item by registration of air nitrogen fluorescence, induced by Extended Air Showers; an effective control over $M \\sim 10^{11}$~ton of atmosphere  may be achieved either from the earth (see Fly's Eye \\cite{FlysEye} and HiRes \\cite{HiRes}) or from satellites (Airwatch \\cite{Airwatch} etc); \\item using such a 'neutrino' detector, as \\emph{Pierre Auger} installation\\footnote{A hybrid detector in Argentina will consist of $1600$ tanks, each filled with $12$~m$^3$ of water. Stations will be distributed in a grid with $1.5$~km spacing. b) Four \\emph{Fly's Eye} type fluorescence detectors, controlling $\\sim 3000$~km$^2$ of the site.} \\cite{Auger}. Note, that \\emph{Pierre Auger} telescope is designed for the study of \\CRs\\ at $E > 1\\times 10^{20}$~eV, high-energy \\n\\ physics being just a byproduct. \\end{itemize} \\subsection{Rate of events} The rate of nucleon-electromagnetic cascades in a \\n\\ detector is proportional to convolution product of $\\nu N$-\\css\\ and $\\nu$-fluxes, and also to the number of scattering centers, i.e.\\ mass of detector. A standard requirement is to have at least $10$ events per year, but the only way to gain higher statistics is to increase the installation mass. But huge arrays are very expensive. However, the higher are \\nuN-\\cs, the higher is the expected rate of events, the better results may be obtained. \\subsection{Background} A severe problem for \\HENA\\ is high \\CR\\ background. One is to shield a detector either burying it deep underground (underwater, underice) or selecting just down-up going events, using the Earth as a shield. In any case, the atmospheric \\n\\ background (mostly from prompt $\\nu$'s) is inevitable. Nevertheless, many $\\nu$-source models suggest that extragalactic diffuse $\\nu$-flux dominates at $E \\gtrsim 10^{15} \\div 10^{17}$~eV. \\subsection{The aims of the paper} We shall discuss several consequences for \\HENA\\ of accelerated growth of \\nN\\ \\css\\ at extremely high energies. We argue that such acceleration may be induced by non-perturbative hard pomeron with intercept $\\sim 1.4$, which was proposed and successfully exploited in series of recent papers by A. Donnachie and P. V. Landshoff \\cite{DL}. We study its influence on the rate of cascades in a $\\nu$-detector and on evolution of \\n\\ spectra during their pass through large column depths of matter. ",
        "conclusions": ""
    },
    "0201/astro-ph0201234_arXiv.txt": {
        "abstract": "{ We present CCD surface photometry of 16 nearby dwarf galaxies, many of which were only recently discovered.  Our sample comprises both isolated galaxies and galaxies that are members of nearby galaxy groups. The observations were obtained in the Johnson B and V bands (and in some cases in Kron-Cousins I). We derive surface brightness profiles, total magnitudes, and integrated colors.  For the 11 galaxies in our sample with distance estimates the absolute B magnitudes lie in the range of $-10 \\ga M_B \\ga -13$.  The central surface brightness ranges from 22.5 to 27.0 mag arcsec$^{-2}$. Most of the dwarf galaxies show exponential light profiles with or without a central light depression. Integrated radial color gradients, where present, appear to indicate a more centrally concentrated younger population and a more extended older population. ",
        "introduction": "Dwarf irregular (dIrr) and dwarf spheroidal (dSph) galaxies account for 80--90\\% of the total population of galaxies.  While there are common trends in their global characteristics such as luminosity, surface brightness, and metallicity, dwarf galaxies may differ significantly in details of their star formation histories and evolutionary state.  Nearby dwarf galaxies out to about 5 Mpc have the advantage that we can study both their integrated properties through medium-sized ground-based telescopes, and their detailed stellar content through high-resolution observations with the Hubble Space Telescope or the new 8m to 10m-class telescopes.  We are carrying out a large project to study both the integrated properties and the resolved stellar content of dwarf galaxies in the Local Volume ($V < 500$ km s$^{-1}$), a necessary precondition for the understanding of the evolution of unresolved dwarf galaxies at larger distances. Hopp \\& Schulte-Ladbeck (\\cite{hs}), Karachentseva et al.\\ (\\cite{karachentseva96}), Bremnes et al. (\\cite{bremnes98};\\cite{bremnes99}), Makarova (\\cite{makarova}), and Jerjen et al. (\\cite{jerjen}) presented the results of surface CCD photometry of many nearby dwarf galaxies within and outside of groups out to a distance of 10 Mpc.  Despite this considerable observational progress, for more than 3/4 of the dwarf galaxies of the Local Volume neither surface brightness profiles nor magnitudes and colors have been measured yet, nor have these galaxies been imaged with modern CCD detectors or reliably classified by morphological type.  Over the last three years Karachentseva and Karachentsev with their co-workers have carried out a search for new nearby dwarf galaxies on the basis of the POSS-II and ESO/SERC sky surveys, covering 97\\% of the sky. Their survey resulted in detection of about 600 dwarf systems more than half of which were missing in the catalogues of known galaxies. Subsequent follow-up observations of these galaxies in the 21 cm hydrogen radio line (Huchtmeier et al.\\ \\cite{huchtmeier}) confirmed that many of these objects are nearby with a median radial velocity of $\\sim$ 1200 km s$^{-1}$. It should be emphasized that during the last two decades the total number of probable Local Volume dwarf galaxies has increased by a factor of two due to new detections and amounts to $\\sim$ 360.  Dwarf galaxy candidates from the latest surveys were included in our ongoing program of ground-based imaging follow-up observations.  For these galaxies we adopt the following naming convention: ``KK'' (Karachentseva \\& Karachentsev \\cite{kk98}), ``KKSG'' (Karachentsev et al.\\ \\cite{k2000}), and ``KKH'' (Karachentsev et al.\\ \\cite{k2001a}), followed by a running number corresponding to the line in each of the respective catalogues.  This sample was supplemented by several low-surface-brightness (LSB) dwarfs located in the M~81 group detected in other surveys. ",
        "conclusions": "It is well known that broad-band color indices are fundamental for studying stellar populations in galaxies, and the color variations along the radius reflect inhomogeneities of the stellar component. About half of the galaxies under investigation demonstrate a minor increase in the redness of the total color index B--V center to periphery. This is likely to correspond to the increase in the average age of the stellar population towards the edge of the galaxy. This property of many dwarf galaxies was noticed earlier (Makarova \\cite{makarova}). In particular, a more extended old stellar population is a common property of dwarf galaxies in the Local Group (Grebel \\cite{grebel}). The absence of a noted color gradient for half of the galaxies of our sample may be indicative of a more homogeneous spatial distribution of the stellar populations of different age in these dwarf galaxies, or of a small age and metallicity spread.  The median color index of the measured galaxies is $\\langle B-V \\rangle = 0\\fm50\\pm0\\fm10$. This color index is somewhat redder than in typical LSB galaxies, where $\\langle B-V \\rangle = 0\\fm45$ (McGaugh \\& Bothun \\cite{mb}; Vennik et al. \\cite{vennik}).  Figure~3 displays the relationship between central surface brightness of the galaxies of the sample discussed and absolute stellar magnitude in the B filter. All the values are corrected for Galactic extinction. For comparison, similar relations for the spiral galaxies from the paper by van der Kruit (\\cite{kruit}), and also for the dwarf galaxies and the galaxies of low surface brightness from the paper by Vennik et al.\\ (\\cite{vennik}) and for the dwarf galaxies from the article by Makarova (\\cite{makarova}) are plotted in the figure. As can be seen from the figure the galaxies of the present study occupy a rather narrow range of absolute B magnitudes, $\\sim -10\\fm0$ to $\\sim -13\\fm2$. Their central surface brightnesses are distributed in the interval from 22 to 25.5 mag arcsec$^{-2}$. The mean value of the central surface brightness in the B filter (corrected for Galactic extinction) for our galaxies is 24.6~$\\pm$~1.3 mag~arcsec$^{-2}$. The nearby dwarf galaxies measured by Vennik et al.\\ (\\cite{vennik}) and Makarova (\\cite{makarova}) have on the average higher absolute magnitudes and central surface brightnesses. The regions of nearby dwarf galaxies and bright spirals are well-separated. Note that there may be a weak correlation of the absolute stellar magnitude and the central surface brightness for the galaxies indicated in the figure. Such a correlation was noted earlier (Binggeli \\& Cameron \\cite{bc}, Vennik et al. \\cite{vennik} and other authors). There is a noticeable separation in the M -- $\\mu$ diagram between galaxies of different morphological types. It was suggested by Binggeli (\\cite{b}) that the M -- $\\mu$ diagram for stellar systems might be the equivalent of the HR diagram for stars (see Fig.~1 of his review).  McGaugh \\& Bothun (\\cite{mb}) believe that the distribution of galaxies in the ($\\mu_0$, M$_B$) plane may reflect the initial conditions of formation of these objects. The luminosity corresponds approximately to the mass enclosed in the density fluctuation from which the galaxy formed, while the surface brightness corresponds to the density gradient in this fluctuation."
    },
    "0201/astro-ph0201144_arXiv.txt": {
        "abstract": "We consider the use of polarization properties as a means to discriminate between Synchrotron Self-Absorption (SSA) and Free--Free Absorption (FFA) in GHz-Peaked Spectrum (GPS) sources.  The polarization position angle (PA) of synchrotron radiation at high frequencies for the optically thin regime is perpendicular to the magnetic field, whereas it is parallel to the magnetic field at low frequencies for the optically thick regime.  Therefore, SSA produces a change in PA of $90^{\\circ}$ across the spectral peak, while FFA does not result in such a change.  We analyzed polarization data from VLA observations for six GPS sources to see if such a change in PA was present.  Our results indicate that there is no significant evidence for $90^{\\circ}$ change in PA across the spectral peak, suggesting that FFA is more likely than SSA for low-frequency cutoffs in these sources. ",
        "introduction": "GHz-Peaked Spectrum (GPS) sources are characterized by a simple convex radio spectrum with a peak at GHz frequencies. The common properties of bright samples of GPS sources are: small size ($<$ 1 kpc), powerful radio luminosity ($L>10^{45}$ erg s$^{-1}$), and apparently low variability (e.g., \\cite{Stanghellini1998}; \\cite{O'Dea1998}). GPS sources associated with quasars tend to show a complex morphology in VLBI images, while those associated with galaxies are likely to be double or triple radio sources (\\cite{Phillip-Mutel1980}; \\cite{O'Dea1991}).  The spectral shape of GPS sources is distinct from that of extended radio sources. Generally, extended radio galaxies have steep spectral profiles, in contrast to the flat spectral profiles of quasars. GPS sources have steep spectra at high frequencies, as do radio galaxies, which indicates optically thin synchrotron radiation. However, the low-frequency cutoff in the spectral profile is peculiar to GPS sources. There has been a considerable debate regarding the interpretation of this low-frequency cutoff.  The low-frequency cutoff is ascribed by some to Synchrotron Self-Absorption (SSA).  SSA is usually discussed in terms of the equipartition brightness temperature, $T_{\\rm B}$ (assuming equipartition of energy between the radiating particles and the magnetic field), which is derived from the flux density and frequency of the spectral peak, and the optically thin spectral index. \\citet{Readhead1996} showed that the observed peak brightness temperature is consistent with the equipartition brightness temperature for the lobes of 2352+495. \\cite{Snellen2000} suggested SSA based on correlations between the peak frequency and the angular size, and between the peak flux density and the angular size.  On the other hand, \\citet{Bicknell1997} suggested that the observed anticorrelation between $\\nu_{\\rm m}$ and size can be explained by Free--Free Absorption (FFA). VSOP and multi-frequency ground VLBI observations of OQ 208 have been used to infer that FFA is the more likely mechanism, with the two radio lobes surrounded by a cold dense plasma which causes thermal FFA (\\cite{Kameno2000}). Assuming intrinsically symmetric double lobes, the apparent asymmetry of opacities towards lobes was explained as being caused by the differences in path lengths through the plasma, which has an estimated electron density and electron temperature of $ 600 < n_{\\rm e} < 7\\times 10^{5}$ and $10^{4} < T_{\\rm e} < 6\\times 10^{7}$, respectively. In addition, FFA is favored for other GPS sources, such as 0108+388 \\citep{Marr2001} and NGC 1052 \\citep{Kameno2001}.  So far, the origin of the low-frequency cutoff has been discussed in terms of the spectral properties and brightness temperatures. However, since both FFA and SSA can produce similar spectral shapes, it is difficult to differentiate between the two mechanisms in this way. Here, we examine the possibility of discriminating FFA from SSA by using the polarization properties of sources. In section~2 we review the relevant polarization properties, and we describe the data analysis in section 3, followed by section 4 where the results for the sources studied to date are given. ",
        "conclusions": "Although five of the six sources have $F_{\\rm{rat}} > 1$, the case for FFA is not compelling, in part because of lack of observed frequencies, particularly around $\\nu_{\\rm m}$. The SSA model predicts not only the $90^{\\circ}$ jump in PA, but also the change in the fractional polarization as a function of the wavelength. In the optically thin regime, the degree of polarization decreases with decreasing frequency, and becomes null at $\\nu_{\\rm m}$ (\\cite{Aller1970}).  It increases again at the optically thick regime, but the polarization degree is $\\sim$ 1/10 of that at the optically thin regime. The percentage polarization is plotted as a function of $\\lambda^2$ for the six sources studied here in figure 5. The variation in the degree of polarization expected from the SSA model cannot be seen for any source except 0738+313. This also implies that SSA is unlikely to be the cause of the low-frequency cutoff.  For 2134+004, $F_{\\rm rat}<$1 results from the low spatial resolution against the complex structure of the source components.  In fact, 2134+004 has two components, which have different $RM$s \\citep{Taylor2000}.  In such a case, our method should be applied to the individual components.  The $RM$ asymmetry of 2134+004 is easily explained by means of the FFA model.  Suppose that two intrinsically identical lobes are surrounded by dense plasma to produce FFA. Because there is usually a difference in the path lengths to the lobes in the plasma, the $RM$ differences reflect the differences in the path length. Such a model allows us to estimate the electron density and the magnetic field \\citep{Kameno2000}.  In conclusion, studies of the polarization properties of six published GPS source have yielded no convincing evidence for the $90^{\\circ}$ change in PA across the spectral peak expected from SSA models. This and the distribution of fractional polarization as a function of wavelength lead us to conclude that FFA may be the dominant cause for the low-frequency cutoffs in most of these sources.  \\bigskip \\bigskip \\bigskip \\bigskip This research made use of data from the University of Michigan Radio Astronomy Observatory, which is supported by funds from the University of Michigan. We thank Philip G. Edwards for much advice."
    },
    "0201/astro-ph0201372_arXiv.txt": {
        "abstract": "{ We have developed a new evolutionary synthesis code, which incorporates the output from chemical evolution models. We compare results of this new code with other published codes, and we apply it to the irregular galaxy NGC 1560 using sophisticated chemical evolution models. The code makes important contributions in two areas: a) the building of synthetic populations with time-dependent star formation rates and stellar populations of different metallicities; b) the extension of the set of stellar tracks from the Geneva group by adding the AGB phases for $m_i/M_\\odot \\geq 0.8$ as well as the very low mass stars. Our code predicts spectra, broad band colors, and Lick indices by using a spectra library, which cover a more complete grid of stellar parameters. The application of the code with the chemical models to the galaxy NGC 1560 constrain the star formation age for its stellar population around 10.0 Gy. ",
        "introduction": "\\label{sec:int}     The study of the history of the evolution of galaxies includes three important issues: {\\it spectral}, {\\it dynamical}, and {\\it chemical evolution}. Since these three different parts of the galaxy evolution are very difficult to study in simple models. Different authors has been focused to one of these parts. Although several efforts have been devoted by different authors (for example \\cite{tinsley76} and \\cite{vazdekis96}) to build complete spectro-chemical evolution codes. It has been difficult to incorporate the new releases in stellar models (structure evolution, stellar yields and stellar atmospheres) to the complexity of galaxy evolution models. The last reason has resulted in the division of the study of the spectro-chemical galaxy evolution in two complementary ways. While the chemical evolution models are the independent variable in the complete spectro-chemical evolution, the spectral evolution codes, which use population synthesis or evolutionary synthesis are left to use simple star formation histories to constrain the results and to be independent models.  In this way, we find sophisticated chemical evolution models. But, the most part of spectral evolution models consider the star formation in a simple star-burst, and a complicated star formation history can be modeled by many star-bursts of different metallicities and star formation rates. Thus, instead of studying a broad range of galaxies using simple star-burst models, it is preferable to develop spectro-chemical evolution models of neighboring galaxies using constrains to model the systems, and after that use the models to galaxies at higher red-shifts.  The new code ({\\code}) presented here has two important contributions to the evolutionary synthesis models: the first one is the consideration of a more complete set of evolutionary tracks from the Geneva group, adding tracks from \\cite{chabrier97} to the very low mass range; the second one is the construction of synthetic populations following the chemical evolution model in order to derive the spectral properties for systems with arbitrary enrichment history.  In this paper we initially use simple star formation scenarios to test the new evolutionary synthesis code, and then we proceed to more complete chemical evolution models to constrain the age of stellar population in the irregular galaxy NGC 1560.  In Section \\ref{sec:trac} we illustrate how the tracks are assembled to built a complete set of tracks for either the high mass loss rates, or the normal mass loss rates, and we describe the spectra library used in our code. In Section \\ref{sec:syn}, we describe the process followed to built the evolutionary synthesis code, and the transformation to the observational plane of the variables. In Section \\ref{sec:models} we test and compare {\\code}  with others codes commonly used in the field and we present a first application by predicting spectral variables for the irregular galaxy NGC 1560. Finally, we discuss our conclusions. ",
        "conclusions": "We have extended the tracks from the Geneva group, producing two sets of tracks, one with high mass loss rate for masses in the range of $12.0 \\leq m_i/M_\\odot \\leq 120.0$ and the other with normal mass loss rate in the range of $0.08 \\leq m_i/M_\\odot \\leq 120.0$. These sets of tracks are complete and consistent for metallicities in the range $0.001 \\leq Z \\leq 0.04$, they cover phases from main sequence up to asymptotic giant branch in the range $0.8 \\leq m_i/M_\\odot \\leq 120.0$, and pre-main sequence and main sequence for masses in $0.08 \\leq m_i/M_\\odot \\leq 0.7$. Together with the spectral library, the code {\\code} is able to predict the spectral properties from stellar populations under a wide variety of conditions.  With this code, it is possible to build synthetic populations as the sophisticated chemical evolution models predict, as long as metallicity is in the range $0.001 \\leq Z \\leq 0.04$.  Results for star-burst evolution in broad band colors, spectral indices and spectra from our code fit very well with the same results obtained with other codes. {\\code} could be used in young stellar populations like in star-burst or blue galaxies, or stellar populations inside cores of AGN's. In the same way, the code could be used for old stellar populations like those in early type galaxies and all galaxy types in the middle of young and old stellar populations.  In applying just chemical models using observational constraints mentioned here for NGC 1560 it is not easy to distinguish what it is the age of the population. But using observed colors for this galaxy a spectro-chemical evolution model can be fitted and a population around 10 Gy and mean metallicity of stars $Z=0.002$ are predicted with our model.  In the same way, from Figure \\ref{f14}, the evolution of colors for the 10.0 Gy model is possible to see that colors match the box error at 5 Gy. Furthermore, we might say that star formation history age for the mean stellar population is between [5.0,10.0 Gy], though we do not have chemical model for 5 Gy.  We have obtained success in reproducing the values for many simple star formation scenarios, and finally, we have matched very closely the values observed for two broad band colors in NGC 1560."
    },
    "0201/astro-ph0201414_arXiv.txt": {
        "abstract": "A catalog containing milliarcsecond--accurate positions of 1332 extragalactic radio sources distributed over the northern sky is presented -- the Very Long Baseline Array Calibrator Survey (VCS1). The positions have been derived from astrometric analysis of dual--frequency 2.3 and 8.4 GHz VLBA snapshot observations; in a majority of cases, images of the sources are also available. These radio sources are suitable for use in geodetic and astrometric experiments, and as phase--reference calibrators in high--sensitivity astronomical imaging. The VCS1 is the largest high--resolution radio survey ever undertaken, and triples the number of sources available to the radio astronomy community for VLBI applications. In addition to the astrometric role, this survey can be used in active galactic nuclei, Galactic, gravitational lens and cosmological studies. The VCS1 catalog is available at \\catref. ",
        "introduction": "Accurate celestial and terrestrial reference frames play an important role in many areas of science.  Celestial reference frames have been used historically for navigation (terrestrial and deep--space), time keeping and for studying the dynamics of solar system, Galactic and extragalactic objects. For over three decades, high--accuracy radio interferometry of extragalactic radio sources using the techniques of Very Long Baseline Interferometry (VLBI) has played a major role in defining these terrestrial and celestial frames \\citep{sovers98}.  Precise measurements of orientation, rotation and deformation of the Earth's surface provides unique information concerning its internal structure, and allows detailed testing of geodynamic theories (e.g. plate tectonics) and studies of the major geophysical fluids (the atmosphere, oceans and groundwater) \\citep{eubanks93}.  The International Earth Rotation Service (IERS -- http://www.iers.org) was established in 1988 by the International Astronomical Union (IAU) and the International Union of Geodesy and Geophysics (IUGG) to serve the astronomical, geodetic and geophysical communities by providing an International Celestial Reference System (ICRS) and an International Terrestrial Reference System (ITRS) to define fundamental inertial frames of reference for celestial and terrestrial positions. The current realizations of the ICRS and ITRS are the International Celestial Reference Frame (ICRF) \\citep{ma98} and the International Terrestrial Reference Frame (ITRF2000) \\citep{altamimi02}. Physical connection of these two systems depends on accurate measurement of two celestial angles (the offset in longitude and obliquity of the celestial pole with respect to its position defined by the conventional IAU precession/nutation models -- $\\Delta\\psi$ and $\\Delta\\epsilon$) and three Earth orientation parameters (EOPs): UT1 (changes in the length of day due to variations in the rotation of the Earth), and the X \\& Y polar motion offsets.  The ICRF was adopted by the IAU as of January 1998 as the realization of the ICRS, replacing the FK5 \\citep{fricke88}. It is a kinematic reference frame, based on the positions of distant quasars and active galactic nuclei, rather than a dynamical reference frame such as the FK5 frame (which incorporates the Earth's motion, the mean equator and the dynamical equinox at some reference epoch). The origin of the ICRF is at the solar system barycenter, and its axes are defined by the directions to a set of extragalactic radio sources measured using VLBI. The ICRF axes are consistent with the J2000 alignment of the FK5 to within the accuracy of the FK5. The orientation of the ICRF is defined by the positions of the 212 defining sources included in the initial ICRF solution \\citep{ma97}, with an estimated accuracy of 0.25 milliarcseconds (mas). A further 396 sources were included in \\citet{ma98} as candidate ICRF sources and to improve the catalog density, and subsequently 59 additional candidate sources were presented in ICRF--Extension 1 (\\icrfext) \\citep{iers99}.  The data used to define the ICRF included over 2.2 million VLBI observations made between 1979--1995. Significant improvement of the current ICRF will require VLBI monitoring of the mas structure of the current sources, and the addition of new sources to the frame solutions.  Astronomical imaging of weak radio sources using phase--referencing also relies on a dense and accurate grid of bright radio sources. Phase referencing \\citep{beasley95} involves rapid ($\\sim$minutes) antenna position--switching between an astronomical target source and an adjacent (1--5\\arcdeg\\, separation) calibrator. By interpolating antenna--based phase, delay and rate corrections derived from the calibrator observations to the weak target source, phase coherence can be extended indefinitely, allowing longer integration times and thermal--noise--limited imaging. To first order, unmodeled antenna--based geometric and electronic delays are removed. Second--order errors in this process increase with (a) switching time between observations of the calibrator, which depends on antenna slew rates, source brightness and system sensitivity, and (b) the angular distance to the calibrator (i.e. a breakdown of isoplanicity).  Dense calibrator grids are required to minimize these errors.  Over the past five years, phase--referencing has allowed imaging of weak radio sources including GRBs \\citep{taylor99}, radio stars \\citep{beasley00} and deep--field radio sources \\citep{garrett01}.  Another advantage of phase referencing is the astrometric registration of multi--epoch observations, such as images of radio supernovae \\citep{bartel00}, maser complexes \\citep{herrnstein99} and the Galactic center \\citep{reid99}.  In this paper we present the results of a multi--epoch dual--frequency Very Long Baseline Array (VLBA) survey -- the VLBA Calibrator Survey (VCS1) -- of 1811 VLBI sources identified on the basis of high--resolution Very Large Array observations. Some 1332 of these sources have not been previously measured in astrometric mode -- their positions are presented here as the VCS1 catalog (see \\catref). Future VLBA calibrator surveys (e.g. VCS2 -- \\citet{fomalont01}) will build on this initial catalog. The goals of the survey were: (a) to increase the surface density of known geodetic--grade calibrators with mas--accurate positions in the northern sky, providing candidate sources for future extensions of the ICRF; (b) to facilitate routine phase--referencing to most regions of the northern and equatorial sky, allowing high--resolution radio imaging of weak scientific targets; and (c) to provide a uniform image database at 2.3 and 8.4 GHz for use in scientific applications, including AGN \\& gravitational lensing studies, and cosmology. Milliarcsecond--accurate positions in the reference frame of the ICRF and a selection of dual--frequency images of VCS1 sources are presented. In Section 2 we describe the VCS1 sample and the survey observations, and present a selection of survey images. In Sections 3 \\& 4 we discuss the astrometric analysis of the VCS1 sample, and examine ongoing efforts to use VCS1 for science and to expand it through the Galactic Plane and to higher frequencies. ",
        "conclusions": "The mas--accurate positions for compact bright extragalactic radio sources presented in this survey will enable phase--referencing VLBI imaging of weak astronomical targets over large areas of the northern sky.  The combined \\icrfext\\, and VCS1 sky density is sufficient to provide a calibrator within 3\\arcdeg\\, of a random location north of -30\\arcdeg\\, declination approximately 75\\% of the time, and within 5\\arcdeg\\, in 96\\% of cases. Experiments requiring high astrometric accuracy, or in cases where no suitable cataloged VLBI calibrator can be found within a few degrees, may use weaker continuum sources close to the astronomical target if sufficient data recording bandwidth is available. These weaker continuum sources may be identified from lower--resolution surveys such as the NRAO VLA Sky Survey \\citep{condon98}.  Short and long--term source variability may vary the flux densities of the VCS1 sources by tens of percent in extreme cases, so weaker calibrators should be examined before use in critical applications.  Use of the VCS1 in scientific studies has already commenced. The parent survey to VCS1 -- JVAS -- has been extensively used to search for gravitationally--lensed systems \\citep{patnaik92b,king99}. The existence or absence of gravitational--lens pairs on mas--scales can be used to place limits on the cosmological abundance of supermassive compact objects in the mass range $\\sim10^6$\\, to $10^8$\\,M$_\\sun$\\, \\citep{wilkinson01}. Current studies based on samples of $\\sim$300 sources indicate that such objects cannot make up more than ~1\\% of the closure density of the universe. This suggests that a population of supermassive black holes forming soon after the Big Bang does not contribute significantly to the dark matter content of the Universe.  Examination of the VCS1 images to identify compact--symmetric objects (CSOs) has been carried out \\citep{peck00a}, doubling the number of these objects available for follow--up observations, including the detection of neutral hydrogen absorption towards the radio components of some CSOs \\citep{peck00b}, which may indicate the presence of an obscuring atomic torus in the nucleus of these objects. CSOs have also been found to be remarkably stable flux calibrators \\citep{fassnacht01}, so larger samples will prove valuable for VLA and VLBA flux monitoring experiments, such as measuring time delays in gravitationally lensed sources.  Additional surveys based on the VCS1 to identify sources suitable as geodetic and phase--referencing calibrators at high frequencies (22--90 GHz), and throughout the Galactic Plane, have recently commenced \\citep{fomalont01}.  Proper--motion studies of Galactic objects such as pulsars \\citep{brisken00} depend on dense grids of suitable low--frequency calibrators at low Galactic declinations, which VCS1 does not cover.  The results of these efforts should further benefit the geodetic and astrometric communities.  At the present time, radio observations are the most accurate way to define the ICRS; however in future optical observations will likely play an important role. The Hipparcos stellar reference frame \\citep{perryman97} has been aligned with the ICRF to within 0.6 mas offset and  0.25 mas in rotation at epoch 1991.25 \\citep{kovalevsky97}, and represents the optical realization of the ICRS. Over the next two decades, new optical interferometers and astrometry missions such as the NASA Space Interferometry Mission \\citep{shao98}, the US Naval Observatory Full--Sky Astrometric Mapping Explorer \\citep{horner00} and the European Space Agency's Global Astrometric Interferometer for Astrophysics \\citep{perryman01} will achieve microarcsecond positional accuracies, requiring new definitions of the ICRS."
    },
    "0201/astro-ph0201078_arXiv.txt": {
        "abstract": "The probabilistic nature of the IMF in stellar systems implies that clusters of the same mass and age do not present the same unique values of their observed parameters. Instead they follow a distribution. We address the study of such distributions in terms of their confidence limits that can be obtained by evolutionary synthesis models. These confidence limits can be understood as the inherent uncertainties of synthesis models. We will compare such confidence limits arising from the discreteness of the number of stars obtained with Monte Carlo simulations with the dispersion resulting from an analytical formalism. We give some examples of the effects on the kinetic energy, V--K, EW(H$\\beta$) and multiwavelength continuum. ",
        "introduction": "In recent years, several efforts have been dedicated to improve our understanding of stellar evolution with more detailed and complete theories; at the same time, more powerful observatories have been developed to test the theory.  However, an intermediate tool is necessary to link these pieces of information when we deal with systems in which only the integrated light of stellar populations (and their by-products, like the emission line spectrum) is available: this tool are synthesis models. Recently, it has been established how the input libraries affect the predictions of synthesis models (see \\opencite{Bru01a} or \\opencite{Car00} as examples). From the theoretical point of view, there are still several open questions in the modelization of stellar clusters by evolutionary synthesis codes. One of the most important ones is related with the conservation of energy and the Fuel Consumption Theorem established by \\inlinecite{RB86} (see also \\opencite{MG01} for the link of chemical with spectrophotometric models). Other questions, related with ``technical'' details in the isochrones computation can be found in \\inlinecite{Cetal01a}. We also refer to the contribution of S. Yi in these proceedings.  In addition to those listed, there is still a source of uncertainty arising from the use of the Initial Mass Function (IMF) and the effect of the discreteness in the number of stars in the models results (\\opencite{Buz89}; \\opencite{SF97}; \\opencite{LM99}; \\opencite{CLC00}; \\opencite{Plu01}; \\opencite{Bru01b}; \\opencite{Cetal01b}).  In this paper we present our current understanding of the dispersion introduced in the results of evolutionary synthesis models by the discreteness of the stellar population for a given IMF. ",
        "conclusions": "We have shown that the effects of fluctuations in the number of stars due to the stochastic nature of the stellar formation process and the discreteness of the stellar populations produce a dispersion in the predictions of evolutionary synthesis models, and that such dispersion may be much larger than the observational errors. The dispersion can be evaluated theoretically, and it can be used as an observable.  The application of this ideas may help to improve the understanding of other astrophysical problems, for example: (i) Is it really necessary a IMF slope different from Salpeter's?  (ii) Is it possible to explain the observed dispersion of chemical abundances by including the IMF fluctuations in chemical evolution models?.  (iii) How much the underlying probability distribution of luminosities affects the corresponding colors?..."
    },
    "0201/astro-ph0201287_arXiv.txt": {
        "abstract": "{ We have used a Monte Carlo radiative transfer code to produce edge-on images of dusty galactic disks, allowing a fraction of the dust to be distributed in clumps. Synthetic images of edge-on galaxies have been constructed for different amounts of dust, distributions of clumps and fractions of dust in clumps, following the formalism of \\citet{BianchiSub1999}. We have also considered models with stellar emission embedded in the clumps. The synthetic images have been fitted with analytical models made with smooth distributions of dust, adopting the procedure developed by \\citet{XilourisSub1998} to fit optical images of real edge-on galaxies. We have compared the parameters determined by the fit with the input parameters of the models. For the clumping distributions adopted in this paper, the neglect of clumping results in underestimating the amount of dust in a galaxy. However, the underestimation is never larger than 40\\%. ",
        "introduction": "Optical extinction by interstellar grains complicates the study of the stellar and dust content of a spiral galaxy. For moderate to high optical depths, information on the intrinsic properties of each of these components is entangled in the observed surface brightness distribution of the galaxy. The usual way of deriving these properties is by comparing the observations with radiative transfer calculations in a dusty galactic model. A certain degree of complexity is needed in the description of the geometric distributions of a galactic model. Radiative transfer models with simple geometries have been proven to provide equivocal results \\citep*{DisneyMNRAS1989}. Furthermore, scattering must be taken into account, given the high albedo observed for Galactic dust \\citep*{GordonApJ1997}.  A radiative transfer code using the Monte Carlo (MC) technique can, in principle, accommodate any stellar and dust distribution and provide an accurate calculation of the surface brightness for a given model \\citep*{WittApJ1992,BianchiApJ1996,DeJongA&A1996b}. However, the MC method is time consuming and thus can not be easily included in fitting procedures. For such procedures an approximate treatment of the radiative transfer is more appropriate.  The first approach of such a procedure was introduced by Kylafis \\& Bahcall (1987, hereafter KB), where vertical profiles of a model galaxy were fitted to observations of NGC~891. In KB's work, scattering is calculated up to the first order and an approximation is introduced for the higher order contributions while the distribution of the dust is assumed to be smooth and exponential along the radial and vertical directions. A similar approximation for the treatment of scattering is adopted by Silva et al. (1998). Other works prefer an exact solution for scattering within simple geometries \\citep*{BruzualApJ1988,CorradiMNRAS1996,XuA&A1995}.  Using the KB method, the radiative transfer equation can be solved for a wide variety of geometries \\citep{ByunApJ1994}. Xilouris et al. (1997;1998) improved the original idea of KB and implemented a technique to fit the surface brightness distribution of edge-on spirals. In edge-on galaxies, the effects of extinction are maximized and it is possible to  separate the stellar and dust components. The method has been successfully applied to a sample of seven edge-on galaxies \\citep{XilourisSub1998} concluding that late-type spiral galaxies have a moderate opacity, with a mean face-on optical depth in the B-band $\\tau_\\mathrm{B}\\approx 0.8$. As in most of the models used for the description of spiral galaxies, \\citet{XilourisSub1998} adopt smooth exponential distributions to describe the stellar and dust disks.  Real galaxies exhibit a wide variety of inhomogeneities, like spiral arms, bars, clumps. While it is possible to include these structures in the solution, when fitting an observed image it is desirable to deal with the simplest possible description, in order to limit the number of model parameters. Complex models can then be used to test the reliability of the description obtained with the simple models. This is done, for example, in \\citet{MisiriotisA&A2000}, where synthetic edge-on images of model galaxies with a spiral structure are fitted with a smooth exponential model. They find that plain exponential distributions provide a good description of the galactic disks. The derivation of the stellar and dust parameters is only slightly affected by the spiral pattern.  In this paper we perform a similar exercise to study the influence of dust clumping on the fit of edge-on galaxies. In most cases, a clumpy medium exhibits higher transparency than a smooth one of the same dust mass. Therefore, the comparison of real images with smooth models may result in an underestimation of the dust content in a galaxy. Here we use the MC method  presented by \\citet[][hereafter BFDA]{BianchiSub1999} to create images of clumpy galactic models. For the parametrisation of BFDA, the maximum difference between clumpy and homogeneous models can be seen in the edge-on case. Monte Carlo models of highly inclined galaxies had been produced also by \\citet{KuchinskiAJ1998}, with a different parametrisation for the clump distribution. With respect to their result, the influence of clumping in the edge-on case is larger for the model adopted here (see the discussion in BFDA). The MC images of BFDA will be fitted with the KB model as in \\citet{MisiriotisA&A2000}, aiming at determining the efficiency of clumping in hiding the dust when a galaxy is seen edge-on.  Sect.~\\ref{fits} describes the method adopted for the comparison and the MC models we have used to produce the synthetic images. A model with exponentially distributed clumps developed for this work is presented in the Appendix.  The results are presented in Sect.~\\ref{results} and a  summary is given in Sect.~\\ref{conclu}. ",
        "conclusions": "\\label{conclu}  We have used a Monte Carlo code for the radiative transfer in a clumpy galactic model to produce synthetic edge-on images of spiral galaxies in the V-band. The images have then been analyzed using the fitting procedure of \\citet{XilourisSub1998}, that uses the method of \\citet{KylafisApJ1987} to compute the radiative transfer in smooth galactic models. We first checked for the consistency of the two independent codes, by analyzing smooth MC images. Despite the differences in the solution of the radiative transfer, it was possible to retrieve the parameters of the MC models within 4\\% of their original value and we confirmed the importance of scattering in the radiative transfer simulations. The dust mass is significantly underestimated when scattering is not included in the fitting routine.  We have then studied the influence of clumping on the derivation of the parameters, when a synthetic image of an edge-on clumpy galactic model is analyzed through a simpler (and more controllable) smooth model. As expected, using a smooth model when fitting images of edge-on galaxies leads to an underestimation of the dust content of the galaxy. For the distribution of clumps adopted in this paper, about 20-30\\% of the dust mass is missed by the fitting, the value depending more on the fraction of dust located in clumps rather than on the details of the clumps distribution or the total dust mass of the model. In addition, clumping alters the surface brightness distribution with respect to that of a smooth exponential model, thus making the fit more difficult. This may explain the overestimation of the intrinsic luminosity in models with clumping.  We stress again here that our work is limited to the edge-on case, for which convergent fits can be obtained with the KB technique. The results may also depend on the clump distribution and on the geometrical parameters adopted here. However, it is interesting to note that other works find a small influence of clumping in edge-on galaxies. \\citet{KuchinskiAJ1998} produced radiative transfer models of disk galaxies including clumping. The distribution of clumps is different from the one adopted here: the clumpy medium is characterised by a costant filling factor throughout the whole disk and by a density contrast with the homogeneous medium \\citep[][ see also BFDA for a comparison with the present models]{WittApJ1996}. Comparing colour gradients across the extinction lane of highly inclined galaxies with the model results, they find that the inferred opacity of the dust disk is largely insensitive to the difference between clumpy and homogeneous dust distributions.  Recent analyses of dust emission at $\\lambda>100\\mu$m suggest that the amount of dust in spiral galaxies is larger than what derived from fits of edge-on surface brightness profiles \\citep{BianchiA&A2000,PopescuA&A2000,MisiriotisA&A2001}. Clumping has been imputed to as a possible cause for this discrepancy: dust in star-forming regions may contribute dominantly to the Far Infrared emission without being seen through its extinction effects. However, the underestimation produced by the clumpy distributions adopted in this paper is always smaller than 40\\%, unable to explain the magnitude of the effect. Different dust/clump distributions may be needed. \\citet{PopescuA&A2000} and \\citet{MisiriotisA&A2001} find that the long wavelength emission observed in two galaxies (namely, NGC 891 and NGC 5907) can be explained by including a second dust disk, at least as massive as the dust disk causing the observed extinction lane and thin enough to escape detection through the technique of surface brightness fitting."
    },
    "0201/astro-ph0201215_arXiv.txt": {
        "abstract": "We investigate the energy release due to the large-scale structure formation and the subsequent transfer of energy from larger to smaller scales. We  calculate the power spectra for the large-scale velocity field and show that the coupling of modes results in a transfer of power predominately from larger to smaller scales. We use the concept of cumulative energy for calculating which energy amount is deposited into the small scales during the cosmological structure evolution. To estimate the contribution due to the gravitational interaction only we perform our investigations by means of dark matter simulations. The global mean of the energy transfer increases with redshift $\\sim (z+1)^{3}$; this can be traced back to the similar evolution of the merging rates of dark matter halos.  The global mean energy transfer can be decomposed into its local contributions, which allows to determine the energy injection per mass into a local volume. The obtained energy injection rates are at least comparable with other energy sources driving the interstellar turbulence as, e.g. by the supernova kinetic feedback.  On that basis  we make the crude assumption that processes causing this energy transfer from large to small scales, e.g. the merging of halos, may contribute substantially to drive the ISM turbulence which may eventually result in star formation on much smaller scales. We propose that the ratio of the local energy injection rate to the energy already stored within small-scale motions is a rough measure for the probability of the local star formation efficiency applicable within cosmological large-scale n-body simulations. ",
        "introduction": "\\label{introduction} During the last decade great success could be achieved in our understanding of the detailed mechanisms for the formation and evolution of cosmic structure. The role of the underlying cosmological models and parameters have been investigated by numerous numerical n-body simulations. On large scales the calculated distribution of matter is in excellent agreement with the observed one, e.g. with the galaxy distribution and the distribution of the intergalactic medium (IGM). The fast enhancement of the available computational power permitted to cover an increasing dynamical range and/or to consider additional processes. In particular hydrodynamical models have been developed very successfully.  The evolution of cosmic matter can roughly be subdivided in terms of scales: On spatial scales larger than the Jeans length gravitation dominates and on smaller scales hydrodynamical processes do. On the other hand for the description of the IGM, e.g., this distinction is insufficient. The IGM is strongly influenced by both, the large scale structure evolution and the feedback of the luminous matter, in particular by the star formation processes:  Supernova explosions sweep out the galactic gas enriched by heavy elements into the IGM changing its chemical composition and thermal state.  Radiation ionizates the IGM in the environment of the galaxies.  In order to obtain an appropriate description of the physical state of the large-scale distributed gas also the amount and the distribution of stars and their back-reaction has to be estimated. However, to link the process of star formation to the large-scale structure evolution is by far out of scope.  Incorporating the stellar feed-back in simulations inevitably needs to connect the star formation rate to available gas parameters. Those parameters can be, e.g. the local density and the local gas temperature. Note, in terms of the considered simulations `local' stands for the average over smallest resolved scales, usually of the order of 1 - 100 kpc. \\citet{schmidt:59} found that in the interstellar gas the star formation rate $\\dot{\\rho}_\\ast$ (SFR) is related to the density $\\rho$ by $\\dot{\\rho}_\\ast \\propto \\rho^n$, where $n$ is adopted to be about 1.5 \\citep[cf.][and references therein]{kennicutt:98}. Although those scales are far below the resolution in large-scale simulations, this empirical relation is often applied. Gas which fulfills certain density and temperature criteria is assumed to form stars according to the above given Schmidt-law \\citep[e.g.][]{yepes:97,springel:00,steinmetz:01,nagamine:01,ascasibar:01}. An alternative approach has been introduced by  \\cite{kauffmann:99} linking semi-analytic galaxy models to bound dark matter halos. Applying such prescriptions for star formation, e.g., permits to calculate the star formation history and the stellar metallicity distribution in the universe. However, the variety of used criteria \\citep[cf.][]{kay:01} indicates that the conjunction is still uncertain.  Knowing the processes on the scale of Molecular Clouds (MCs) which most probably are controlling the star formation rate would provide an indication on the possible linking quantities. Therefore, let us shortly summarize some recent results of detailed investigations related to star formation.  Star formation is hosted by interstellar clouds of molecular hydrogen. Stars probably arise from shock-compressed dense cores within the clouds \\citep{blitz:99}. The cores are produced by supersonic motions of the gas due to the presence of turbulence \\citep{burkert:01}.  According to \\citet{klessen:00b} the SFR especially depends on the scales on which turbulence is driven.  There are indications that the formation of clouds is linked to larger scales:  \\cite*{blitz:99} argued that MCs are formed through the condensation of \\Ion{H}{I} regions in conjunction with some other mechanism as has been proposed by \\cite{ballesteros-paredes:99}, e.g. due to  colliding interstellar gas streams \\citep[cf.][]{burkert:01}.  The life-time of the MCs, $20 - 100\\EMyr$, indicates that the internal turbulence is more or less continuously driven by external forces \\citep{maclow:98,padoan:01}.  Nevertheless, star formation occurs probably only once within one cloud \\citep{elmegreen:00}.  Consequently, the rate of star formation may depend on the presence of an effective intercloud turbulence.  The state of the ISM, i.e. whether it is turbulent or not, may be derived from the power spectrum of the velocity field or from that of the matter distribution. For instance, \\citet{goldman:00} investigated the neutral hydrogen distribution of the Small Magellanic Cloud (SMC) and derived the spatial power spectrum.  He found that the spectrum of the SMC is to some extent steeper than Kolmogorov's 1/3-law and explained that by the compressibility of the gas.  Furthermore, he suggested that the origin of the turbulence could be caused by a close encounter between the Small and the Large Magellanic Cloud (LMC) $2\\times10^8\\Eyr$ ago and he showed that supernova feedback could only contribute little to the kinetic energy of the SMC. This illustrates one possible mechanism to force the ISM into a turbulent state: kinetic energy is transferred from larger to smaller scales, i.e.,  kinetic energy from the relative motion of the SMC and LMC is transferred into internal motions.  It is intriguing in respect to the relation between turbulence and star formation that both the SMC and the LMC show evidence for large increases of their star formation rate at the time of the encounter \\citep{larson:01}.  From these investigations one may conclude that the SFR is controlled by the rate of formation of dense substructures - or clouds - in the interstellar medium, dense enough to undergo gravitational collapse and fragmentation. Whether these substructures are generated by thermal instability, gravitational instability, turbulent compression or some other process is still under debate \\citep[cf. introduction in][]{scalo:02}. We suppose that interstellar turbulence is the leading mechanism. Then the question arises, how is the interstellar medium forced into a turbulent state? What drives the turbulence over a sufficient long time period needed to form eventually stars? Again, different energy sources such as supernova feedback, galactic differential rotation \\citep{sellwood:99}, infall of  high-velocity clouds \\citep{blitz:99b} and other mechanism are conceivable. Merging events between galaxies or minor-merging events between galaxies and smaller clouds are supposed to cause efficient star formation \\citep{kolatt:99,somerville:01,kauffmann:01}. This may indicate that merging, which occurs due to the structure formation, could provide sufficient energy to stir the interstellar medium.  Therefore, the question arises how much energy is eventually available by injection from the extra-galactic scales. When large-scale structures are formed gravitational energy is released and stored in large-scale motions. The structure evolution transfers the kinetic energy from large-scales to galactic scales. May this transferred energy significantly contribute for balancing the dissipation of the turbulent field on the sub-galactic scales?  In this paper we want to address the question which energy transfer down to galactic scales can be expected from the large-scale structure formation, what is the time evolution of the transfer and how is it spatially distributed. On large scales ($\\gtrsim 0.5\\EMpc$) the baryonic matter is still tightly coupled to the dark matter, thus the energy transfer is governed by the gravitational interaction only. We use cosmological n-body simulations to determine the transfer. We start with the investigation of the power spectrum of the large-scale velocity field, closely related to the energy spectrum (Sec.~\\ref{sec-vel}). Then by introducing the concept of cumulative energy we determine the energy injection into galactic-scale motions (Sec.~\\ref{sec-culme}). We calculate the volume averaged energy transfer and discuss its relation to the merging rate evolution of halos (Sec.~\\ref{sec-pik}). Furthermore, the spatial distribution of the energy injection at given scale is determined. This allows to infer a local energy injection rate which can be attributed to a heating rate with respect to the baryonic matter (Sec.~\\ref{sec-dsg}).  Making the crude simplification that the obtained energy injection rates at galactic scales are linked to the necessary energy input for driving the interstellar turbulence and moreover assuming that this turbulence is controlling the star formation leads us to an estimator for the local - in terms of cosmological simulations - star formation rate (Sec.~\\ref{sec-sfr}). Finally we summarize our results and discuss possible implications (Sec.~\\ref{summary}). ",
        "conclusions": "\\label{summary}   During the formation of the large-scale structures gravitational energy is transferred into the internal motion of the dark matter halos. In particular by merging of the halos part of the kinetic energy is subsequently transferred from the large-scale movement into small-scale internal motions. The baryonic matter possesses the same amount of kinetic energy per mass as long as it moves together with the dark matter, i.e. as long as gravitational interaction dominates over pressure forces. Thus the large-scale structure formation is transferring energy into small scales. We have addressed the question how this transfer can be described, how it evolves during the cosmological structure formation and whether this transferred energy rate is of comparable order of magnitude to compete with internal energy sources in the galaxies.  We have performed cosmological n-body simulations to determine which energy amount the large-scale structure formation provides to a given small scale.  A crucial point in that picture is how the energy transfer is directed. We have considered this question investigating the evolution of the power spectrum of the large-scale velocity field. To this end we have used a method introduced by \\citet{lomb:76}. The power spectrum evolves according to the linear growth at scales larger $\\approx 60\\,h^{-1}\\EMpc$. Below these scales it possesses a more shallow region and at scales below $\\approx 0.5\\,h^{-1}\\EMpc$ it roughly behaves like a power law $\\sim k^{-4.2}$. In general, the features of the velocity spectrum are similar to those of the power spectrum of the density field. The behavior of the velocity spectra shows that large-scale modes dominate the power of small-scale modes and that the energy transfer throughout the modes is almost exclusively one-directional, namely from larger to smaller scales.  To calculate the energy transfer through a given scale we have used the concept of cumulative energy. The Fourier transforms of the velocity field and of the density of momentum are subdivided into a high-pass and a low-pass filtered part. As a result the energy density can also be decomposed into a high-pass and a low-pass filtered part and the scale-by-scale energy budget contains a term describing the energy exchange between the two spectral ranges. The exchange rate may be calculated by help of the time derivative of the high-pass filtered energy and equals to the energy injection rate into scales smaller than the given cut-off length.  Using the involved quantities we can define straightforwardly the mean global energy exchange rate into given scales at any time during the performed n-body simulations. We obtain the mean energy exchange rate which increases with redshift according to a power law $\\sim (z+1)^{3}$. We argue that this can be attributed to the mass increase of dark matter halos. The evolution includes all kinds of infall, i.e.  continuous accretion of matter onto halos, minor mergers, or even merging of halos. If all processes evolve similar or if merging is dominant the evolution of the mean energy exchange exhibits the same time-dependence as the merging rate of dark matter halos.  Due to the additivity of the cumulative energy with respect to its local contributions if applied to grid based n-body systems a {\\em local} high-pass filtered mean energy exchange rate can be defined supposing that the size of the local averaging volume is equal to the cut-off scale $\\lambda_K$. A specific energy exchange rate is given by the ratio of the local energy exchange rate into a given volume to the local mass density. Using this definition we obtain the mean specific energy exchange rate which is as low as $\\approx 10^{-28}\\,n\\Eerg\\Epers\\Eperccm$. However, its expectation value is proportional to the density. Thus, in regions as dense as the galactic halo the energy input due to the energy exchange may be as large as $\\approx 10^{-24}\\,n\\Eerg\\Epers\\Eperccm$, which is comparable with the kinetic energy feedback by supernova averaged over the same galactic volume.  On the basis of our results we consider a speculative picture: We assume  that the injected energy propagates to MC scales and is therefore available for driving the MC turbulence leading to star formation. On the ground of this quite heuristic assumptions we propose an estimator for the star formation rate in a local volume in cosmological simulations: $\\dot{m}_\\ast \\propto (\\pi_K^> / e_K^>) \\: m$.  The model reproduces approximately the observed global evolution of the star formation rate. Adopting this model means that merging and matter infall processes are considered to enhance the star formation at least or to be the leading processes.     \\newcommand{\\nucphys}{Nucl. Phys.} \\newcommand{\\PRL}{Phys. Rev. Lett.} \\newcommand{\\TUB}{Technische Universit{\\\"a}t Berlin} \\newcommand{\\ZfPhys}{Z. Phys.}"
    },
    "0201/astro-ph0201023_arXiv.txt": {
        "abstract": "{We have constructed an artificial meteor database resembling in all details the real sample collected by the observers of the {\\sl Comets and Meteors Workshop} in the years 1996-1999. The artificial database includes the sporadic meteors and also events from the following showers: Perseids, Aquarid complex, $\\alpha$-Capricornids, July Pegasids and Sagittarids. This database was searched for the presence of the radiants of two weak showers: $\\alpha$-Cygnids and Delphinids. The lack of these radiants in the artificial database and their existence in the real observations suggests that $\\alpha$-Cygnids and Delphinids are the real showers and their radiants could not be formed as an effect of intersections of back prolongated paths of meteors belonging to other showers. ",
        "introduction": "Recently, Polish observers taking part in the {\\sl Comets and Meteors Workshop (CMW)} have reported the rediscovery of two July meteor showers - $\\alpha$-Cygnids and Delphinids (Olech et al. 1999a, 1999b, Stelmach \\& Olech 2000, Wi\\'sniewski \\& Olech 2000, 2001). Both of these showers are weak with maximum Zenithal Hourly Rates (ZHRs) slightly exceeding or laying beneath the sporadic background.  The $\\alpha$-Cygnids are active from the end of June until the end of July. The highest activity with ${\\rm ZHR}=2.4\\pm0.1$ is observed at a solar longitude $\\lambda_\\odot=114.8\\pm0.5^\\circ$. The radiant of the shower at this moment is placed at $\\alpha=305^\\circ$ and $\\delta=+45^\\circ$.  The activity period of the Delphinids is still quite uncertain with the first meteors from this shower detected around July 10 and the last ones as late as the middle of August. According to the recent work of Wi\\'sniewski \\& Olech (2001) the maximum hourly rates are observed at $\\lambda_\\odot=125.0\\pm0.1^\\circ$. The activity at this moment is equal to ${\\rm ZHR}=2.2\\pm0.2$ and the radiant of the shower has the equatorial coordinates equal to $\\alpha=312^\\circ$ and $\\delta=+12^\\circ$.  The equatorial coordinates of the beginnings and ends of meteor paths and its angular velocities for both showers were carefully analyzed using the {\\sc radiant} software (Arlt 1992). This software takes into account the properties of the observed meteors and computes the maps of probability for the presence of a radiant (hereafter PPR maps).  Although PPR maps computed for both of these showers showed distinct features, the resulting radiants were polluted by the influence of the meteors from other showers. A quite strong tail reaching the radiant of the Perseids was detected in the case of the $\\alpha$-Cygnids. There is also a trace of the weak $o$-Draconids radiant in the close vicinity of the radiant of the $\\alpha$-Cygnids (Olech et al, in preparation).  An even more complicated situation is present in the case of the Delphinids. The radiant of this shower is placed not far from the series of ecliptic radiants of the Aquarids complex, $\\alpha$-Capricornids and the Sagittarids.  The radiants of these showers are large and have a complex structure often showing several maxima of activity.  Thus one can suspect that both the $\\alpha$-Cygnids and Delphinids are not the real showers and their radiants produced by {\\sc radiant} software come from crossing the back-prolongated paths of the meteors from other showers active in July and also from sporadic events.  To check this possibility, we decided to construct a realistic database of artificial meteors which thoroughly resembled the real sample analyzed in the above mentioned papers. ",
        "conclusions": "We have compared two samples containing meteors observed in July in years 1996-1999. In the real sample, obtained from real visual observations made by Polish amateur astronomers, we detected clear radiants of $\\alpha$-Cygnid and Delphinid showers. The question, which we wanted to answer was whether these radiants can be produced as the intersections of paths of meteors radiating from the real showers active in July. Thus we constructed the artificial sample resembling in all details the real observations and we included all meteor showers except the $\\alpha$-Cygnids and the Delphinids. These radiants, assuming that they are artificial formations created by intersections of meteors from real showers, should also be seen in the simulated sample. A comparison of both databases showed that it is very difficult to produce circular and clear radiants of the $\\alpha$-Cygnids and the Delphinids using the meteors from an artificial sample. On the other hand such radiants are easy detected in the real sample. This strongly supports the hypothesis that the $\\alpha$-Cygnids and the Delphinids indeed exist.  Finally, we decided to perform another test. To the artificial sample we added the meteors from the $\\alpha$-Cygnids and the Delphinids. These showers were described by the parameters ${\\rm ZHR}_{\\rm max}$, $B$ and $\\lambda_{\\odot (max)}$ given by Stelmach \\& Olech (2000) and Wi\\'sniewski \\& Olech (2001). The numbers of our artificial sample were then increased by 2035 $\\alpha$-Cygnids and 704 Delphinids.   \\begin{figure}[h] \\centering \\includegraphics[scale=.59]{Fig17.ps} \\caption{The PPR map for an artificial sample of 24341 meteors with the $\\alpha$-Cygnids and the Delphinids included and computed for the following parameters: $\\lambda_\\odot=115^\\circ$, $\\Delta\\lambda=1.0^\\circ$ and $V_\\infty=41$ km/s. The maximum distance of the meteor from the radiant is $70^\\circ$. } \\label{FigVibStab} \\end{figure}   \\begin{figure}[h] \\centering \\includegraphics[scale=.59]{Fig18.ps} \\caption{The PPR map for an artificial sample of 23076 meteors with the $\\alpha$-Cygnids and the Delphinids included and computed for the following parameters: $\\lambda_\\odot=125^\\circ$, $\\Delta\\lambda=1.0^\\circ$ and $V_\\infty=35$ km/s. The maximum distance of the meteor from the radiant is $70^\\circ$. } \\label{FigVibStab} \\end{figure}  We calculated the PPR maps for this new database and they are presented in Figs. 17 and 18 (for simplicity we decided to compute the maps for meteors within 70 degrees from the radiant only). Assuming that there are no other showers in July than the Delphinids, $\\alpha$-Cygnids and these listed in Table 1 we expect that our artificial sample should produce the same PPR maps as the real sample (shown in Figs. 9 and 13).  In the case of the $\\alpha$-Cygnids we see that the artificial radiant is more compact than the one obtained from the real sample. However we should expect that there are few poorly known or even unknown showers which are present in the real sample and which were not included in the artificial sample. A good example is the $o$-Draconids shower, which is not listed in the IMO Working List of Meteor Showers, and as it is clearly visible from Fig. 9 is detected in our visual data causing a strong disturbance into the shape of the $\\alpha$-Cygnids radiant.  In the case of the Delphinids both Fig. 18 and especially lower panel of Fig. 13 are similar. The radiant of the Delphinid shower is elongated toward the $\\alpha$-Cygnids radiant. Also the influence of the Aquarid complex is present in both cases.  The similarity of the PPR maps obtained from the new artificial and real samples is another argument for the existence of the $\\alpha$-Cygnids and the Delphinids.  Our artificial databases are accessible via Internet and can be downloaded from the following URL: http://www.astrouw.edu.pl/$\\sim$olech/SIM/. Detailed information about these databases are included in the README file."
    },
    "0201/astro-ph0201509_arXiv.txt": {
        "abstract": "The number of neutrino induced upward going muons from a single gamma ray burst (GRB) expected to be detected by the proposed kilometer scale IceCube detector at the South Pole location has been calculated. The effects of the Lorentz factor, total energy of the GRB emitted in neutrinos and its distance from the observer (redshift)  on the number of neutrino events from that GRB have been examined. The present investigation reveals that there is possibility of exploring physical processes of the early Universe with the proposed kilometer scale IceCube neutrino telescope. ",
        "introduction": "We consider the relativistically expanding fireball model of gamma ray bursts (GRB) \\cite{Eli} and calculate the number of secondary muons from high energy neutrinos of a GRB expected to be detected in a kilometer scale neutrino telescope. In a relativistic expanding fireball protons may be accelerated by Fermi mechanism to energies as high as  $10^{20} eV$. The accelerated protons interact with photons of energies of the order of MeV to produce pions. From protons both $\\pi^{+}$'s and $\\pi^{0}$'s are produced sharing the proton energies roughly in equal parts. The $\\pi^{+}$'s would lead to generation of high energy neutrinos and from $\\pi^{0}$'s high energy photons would be produced.  The decay process for neutrino production from $\\pi^{+}$'s is as follows $\\pi^{+}\\rightarrow\\mu^{+}+\\nu_{\\mu}\\rightarrow e^{+}+\\nu_{e}+\\bar\\nu_{\\mu}+\\nu_{\\mu}$. The Lorentz factor $\\Gamma$ of a GRB has a crucial role in the neutrino production mechanism of the GRB. The efficiency for producing pions in $p-\\gamma$ collisions in the fireball varies as $\\Gamma^{-4}$ and the break energy for neutrino energy spectra varies as $\\Gamma^{2}$. Although $\\Gamma$ grows linearly with the radius in a relativistic fireball model, $\\Gamma$ saturates at a value of the order of 100 \\cite{Eli}.  \\para The fraction of energy lost by protons to pions $f_{\\pi}$ can be expressed as a function of the pair production optical depth $\\tau_{{\\gamma}{\\gamma}}$. If the optical depth is large there would be more pion productions. Also if the optical depth is large then the photons produced from $\\pi^{0}$ decays can not come out as a burst. In that case the GRB would be opaque for photons. The GRBs which are not bright in photons are bright in neutrinos.  \\para The effects of varying  the Lorentz factor $\\Gamma$, the total energy emitted or released in neutrinos $E_{GRB}$ and redshift of the GRB $z$ on the expected number of neutrino induced muons from a GRB have been studied in this work. Each gamma ray burst is an individual phenomenon in the Universe distinguished by its redshift, total energy emitted in neutrinos or photons, duration of the burst and Lorentz factor at a particular instant of time . To obtain informations on the physical process responsible for a gamma ray burst it is always meaningful to carry out studies on the detectability of individual GRBs.  \\para  In Ref. \\cite{Guetta} the correlations of the parameters minimum Lorentz factor $\\Gamma$, wind variability time $t_{v}$, observed photon spectral break energy $E_{\\gamma, MeV}^{b}$ for wind luminosity $L_{w}=10^{53}$ erg are given in FIG.1.for equipartition parameters $\\epsilon_{B}=0.01$ and $\\epsilon_{B}=0.1$. We have taken points from the contour plots given in FIG.1. of Ref. \\cite{Guetta} and used those values in our calculations. In the present work we use the wind luminosity $L_{w}$ but by $E_{GRB}$ we mean the total energy released in neutrino emission. The observed photon spectral break energy $E_{{\\gamma},MeV}^{b}$ is the energy where luminosity per logarithmic photon energy interval peaks.  The fraction of fireball energy that goes into neutrino production is weakly dependent on the wind model parameters. There are two reasons for that. The first reason is for low values of Lorentz factor $\\Gamma$ and wind variability time $t_{v}$ only a small fraction of the pion's energy is converted to neutrinos at high proton energy due to pion and muon synchrotron losses. The second reason is the observational constraints imposed by $\\gamma$-ray observations imply that the wind model parameters $\\Gamma$, wind luminosity $L_{w}$ and wind variability time $t_{v}$ are correlated. We keep in mind the luminosity mentioned in \\cite{Guetta} is the wind luminosity and in this work also we use the same notation. To avoid confusion we mention that the fraction of proton energy that goes into neutrino production does not appear explicitly in the expression of total number of neutrinos emitted from a GRB in our procedure of calculation. The reason behind this is we have used $E_{GRB}$ which is the total energy released in neutrinos to calculate the number of neutrinos emitted from a GRB.   \\para From one proton two muon neutrinos are expected to be generated if both the pion and the muon decay. The neutrinos carry about $5\\%$ of the proton energy, it is given in \\cite{Guetta}. So the total energy $E_{GRB}$ released in neutrinos of a GRB is about $10\\%$ of the original total fireball proton energy. We have used the total energy emitted in neutrinos $E_{GRB}$ to calculate the neutrino spectrum from a GRB. So the energy loss by protons in pion productions as well as the pion and muon energy losses which are the intermediate processes for neutrino productions from protons have been already taken into account in our procedure of calculations. The small value of $E_{GRB}$ compared to the original total fireball energy accounts for the intermediate energy loss processes in production of neutrinos from protons in a GRB. If we use $E_{GRB}=10^{53}$erg in our calculation that means the total fireball proton energy was $10^{54}$erg assuming $10\\%$ of the fireball energy has been released in neutrino emission from the GRB. The wind durations are of the order of 10 seconds.  \\para The neutrino spectrum from a GRB on Earth depends on the total energy emitted in neutrinos $E_{GRB}$, distance of the GRB from the observer $z$ and the neutrino break energy  $E_{\\nu}^{b}$ which is again  dependent on the Lorentz factor $\\Gamma$, observed photon spectral break energy $E_{{\\gamma},MeV}^{b}$, the wind variability time $t_{v}$ as well as wind luminosity $L_{w}$ and equipartition parameter $\\epsilon_{B}$.  \\para Earlier average neutrino event rates in a neutrino telescope of $km^{2}$ area per year from gamma ray bursts have been obtained considering burst-to-burst fluctuations in distance and energy \\cite{hal,alv}. They have used energy and redshift distribution functions of the GRBs  modeled according to observations on GRBs. The neutrino event rates per year from GRBs have been obtained including the fluctuations in the Lorentz factor $\\Gamma$. The purpose of the present work is not to calculate the neutrino event rate expected per year in a muon detector of $km^{2}$ area. We investigate for what range of values of their physical parameters ($E_{GRB}$, $z$, $\\Gamma$)  individual GRBs will be detectable by a kilometer scale neutrino telescope of a given threshold energy.  \\para In GRBs $\\nu_{\\mu}+\\bar\\nu_{\\mu}$'s are almost $10^{5}$ times more abundant than $\\nu_{\\tau}+\\bar\\nu_{\\tau}$. The primary neutrino flux from GRBs contain very small number of $\\nu_{\\tau}+\\bar\\nu_{\\tau}$. Due to neutrino oscillations some of the $\\nu_{\\mu}$ 's convert to $\\nu_{\\tau}$'s during their propagation from the burst location to the detector on Earth . According to SuperKamiokande measurements \\cite{sup} the flux ratio of muon neutrinos and tau neutrinos is  $F_{\\nu_{\\tau}/\\nu_{\\mu}}=0.5$. The $\\nu_{\\tau}+\\bar\\nu_{\\tau}$ events would produce different experimental signatures from the $\\nu_{\\mu}+\\bar\\nu_{\\mu}$ events \\cite{alv}. Cosmic tau neutrinos while coming close to the surface of the detector may undergo a charged current deep inelastic scattering with nuclei inside or near the detector and produce a tau lepton in addition to a hadronic shower. This tau lepton traverses a distance, on average proportional to its energy, before it decays back into a tau neutrino and a second shower most often induced by decaying hadrons. The second shower is expected to carry about twice as much energy as the first and such double shower signals are commonly referred to as double bangs. The tau leptons produced in this way are not expected to have any other relevant interactions as they are losing energy very fast and subsequently they decay to muons. Other than double bang events there are $\\nu_{\\tau}+\\bar\\nu_{\\tau}$ events in which muons would be detected from $\\tau$ decays $(\\nu_{\\tau}\\rightarrow \\tau\\rightarrow\\mu$). The double bang events are much less compared to the events  in which tau leptons decay to muons ($\\nu_{\\tau}\\rightarrow\\tau\\rightarrow\\mu$). The $\\nu_{\\mu}+\\bar\\nu_{\\mu}$ events are much more in number compared to the $\\nu_{\\tau}+\\bar\\nu_{\\tau}$ events. The  relative magnitudes of rates of double bang events from $\\nu_{\\tau}$'s, events in which $\\tau$'s decay to $\\mu$'s and the $\\nu_{\\mu}+\\bar\\nu_{\\mu}$ events per year in a $km^{2}$ area neutrino telescope are given in \\cite{alv}. The majority of muon events in the $km^{2}$ area muon detector would be from interactions of $\\nu_{\\mu}+\\bar\\nu_{\\mu}$ in rock (in case of upward going muons) or ice (in case of downward going muons) at the South Pole location.  In this work we consider only the upward going muon events from $\\nu_{\\mu}+\\bar\\nu_{\\mu}$ of a gamma ray burst because incase of upward going events the background noise is comparatively less.  \\para Recent observations on GRB afterglow have confirmed the relativistic fireball model, to be more specific the internal-external shocks model \\cite{piran}. If GRB neutrinos could be observed by neutrino telescopes it would be a great achievement for understanding the distribution in Lorentz factor $\\Gamma$, the ultrahigh energy neutrino background, testing special relativity and the equivalence principle.  \\para In section 2 we discuss about the parametrization of $\\nu_{\\mu}+\\bar\\nu_{\\mu}$ spectrum from a GRB and the procedure used here to obtain the number of neutrino induced muon signals. In section 3 we mention about the atmospheric neutrino background which often makes the detection of neutrino signals from astrophysical sources difficult in muon detectors. In section 4 the results of this work have been discussed. ",
        "conclusions": "We have investigated on the capability of the next generation kilometer scale IceCube detector in detecting high energy neutrinos from individual GRBs. It is always easier to detect nearby GRBs. Although the neutrino events steeply fall with increasing redshift of the GRB there is possibility of neutrino detection even from the far away bursts in IceCube neutrino telescope. If the GRB is of short duration $t_{\\nu}$, low Lorentz factor $\\Gamma$, emit large amount of energy $E_{GRB}$ in neutrino emission and if its photon  spectral break energy $E_{{\\gamma}, MeV}^{b}$ is high the possibility of detecting neutrinos from that GRB increases. Not only understanding the physics of GRBs, testing special relativity and equivalence principle but we can also look forward to study physical processes occuring in the early Universe with neutrino telescope of $km^{2}$ area."
    },
    "0201/astro-ph0201353_arXiv.txt": {
        "abstract": "We present complementary data on 5 intermediate redshift ($0.44 \\le z \\le 0.66$) Mg~II absorbing galaxies, combining high spatial resolution imaging from {\\it Hubble Space Telescope}, high--resolution QSO spectroscopy from Keck/HIRES, and galaxy kinematics from intermediate resolution spectroscopy using Keck/LRIS. These data allow a direct comparison of the kinematics of gas at large galactocentric impact parameters with the galaxy kinematics obtained from the faint galaxy spectroscopy. All 5 galaxies appear to be relatively normal spirals, with measured rotation curves yielding circular velocities in the range $100 \\le v_c \\le 260$ \\kms. The QSO sightlines have projected impact parameters to the absorbing galaxies in the range $14.5{\\rm h}^{-1} \\le d \\le 75{\\rm h}^{-1}$ kpc; the galaxies have inclination angles with respect to the line of sight ranging from 40 to 75 deg. We find that in 4 of the 5 cases examined, the velocities of all of the Mg~II absorption components lie entirely to one side of the galaxy systemic redshift. The fifth case, for which the galaxy is much less luminous than the other 4, has narrow absorption centered at zero velocity with respect to systemic despite having the largest disk inclination angle in the sample.  These observations are consistent with rotation being dominant for the absorbing gas kinematics; however, the total {\\it range} of velocities observed is inconsistent with simple disk rotation in every case. Simple kinematic models that simultaneously explain both the systemic offset of the absorbing material relative to the galaxy redshifts, {\\it and} the total velocity width spanned by the absorption, require either extremely thick rotating gas layers, rotation velocities that vary with $z$ height above the extrapolation of the galactic plane, or both. In any case, our small sample suggests that rotating ``halo'' gas is a common feature of intermediate redshift spiral galaxies, and that the kinematic signature of rotation dominates over radial infall or outflow even for gas well away from the galactic plane. We discuss possible explanations for this behavior, and compare our observations to possible local analogs. ",
        "introduction": "Metallic absorption lines in the spectra of background QSOs and the nature of their association with galaxies have been subjects of considerable interest over the last decade. A number of circumstantial pieces of evidence, accumulated during the first $\\sim 20$ years of QSO absorption line research, pointed to galaxies as being responsible: e.g. the tendency for systems to split into complexes of total velocity extent consistent with galaxy-sized potential wells, the presence of metals which were presumably produced {\\it in situ}, and clustering properties that resembled that of galaxies. However, exploring the details of the connection between absorption systems and galaxies did not become possible until later work began to directly identify galaxies responsible for individual QSO absorption line systems (Bergeron \\& Boiss\\'e 1991; Steidel, Dickinson,\\& Persson 1994 (SDP); Le Brun \\et 1997). This connection is extremely interesting in the context of our understanding of galaxy evolution and galactic structure because QSO absorption line systems have the potential to explore the gas-phase physical conditions, geometry, and kinematics of galaxies with what can be many orders of magnitude increased sensitivity compared to direct observations of the galaxies. Moreover, the sensitivity of absorption line measurements is largely independent of galaxy redshift, and thus provides the opportunity to study the evolution, over very large time baselines, of characteristics that can be much more subtle than those accessible using traditional faint galaxy techniques.  At present, most identification of absorption systems with individual galaxies is tenuous at best. If a faint galaxy is found near the line of sight with a redshift in agreement with the measured redshift of the QSO absorption system, it is generally taken as as positive identification of the actual object producing the absorption. This type of statistical approach has been adopted by most studies of intermediate redshift metallic absorption systems (e.g., BB91; SDP; Le Brun \\et 1997). There are very few cases where relatively exhaustive spectroscopy down to faint magnitude limits has been obtained for all objects within several hundred kpc of the QSO line of sight (cf. Steidel \\et 1997), making the identifications of each absorber more secure. Possible selection effects inherent in this type of identification procedure have been outlined by Charlton \\& Churchill (1996).  The state of knowledge of the types of galaxies producing various classes of QSO absorption systems varies considerably as a function of redshift and of absorption system taxonomy. The most work has been done at $z \\simlt 1$, for absorption line systems selected by the presence of Mg~II $\\lambda\\lambda 2796$, 2803 doublets with absorption line rest-frame equivalent widths $W_{\\lambda} > 0.3$ \\AA\\ (BB91; SDP, Steidel 1995). These relatively strong Mg~II--selected systems are generally associated with gas having N(H~I)$ \\simgt 10^{17}$ cm$^{-2}$ (i.e., ``Lyman limit systems''), and so would be expected to probe (on average) the outer parts of galaxies where the H~I is more highly ionized than disk gas observed for nearby galaxies. The galaxies responsible for the Mg~II absorption at intermediate redshift ($\\langle z \\rangle \\simeq 0.6$), statistically, appear to be drawn from normal field galaxies with luminosities within $\\sim 1.5$ magnitude of present--day $L^{\\ast}$ (BB91; SDP94). More recent morphological studies using {\\it HST}, of which the data in this paper are a subset, have shown that the identified galaxies are generally of relatively normal morphologies identifiable along the Hubble sequence (Dickinson \\& Steidel 1996; Steidel 1998, Steidel \\et 1997). The colors and magnitudes of these galaxies appear to exhibit little or no evolution with redshift over the range $0.3 \\simlt z \\simlt 0.9$, in agreement with general studies of field galaxy evolution (e.g., SDP, Lilly \\et 1996; Vogt \\et 1996).  There is some evidence that the observed distribution of ``impact parameters'', the projected physical distance between the putative absorbing galaxy and the QSO sightline at the galaxy redshift, is consistent with a roughly spherical distribution of gas of radius $R(L)$ that is a weak function of luminosity (SDP, Steidel 1995, 1998): $$R(L_K) \\simeq 38h^{-1} kpc \\left(L_K \\over L^{\\ast}_K\\right)^{0.2}.$$ A flattened disk--like geometry is not as favored by the existing statistics, but any geometry in which the gas layer has finite thickness compared to its radial extent is very difficult to rule out with present statistics on impact parameters and the incidence of interlopers (cf. Charlton \\& Churchill 1996). Given the available statistics, however, and the likely stochastic nature of the sizes and shapes of gaseous envelopes around field galaxies, simple geometric pictures should be viewed as no more than working models.  Damped Lyman $\\alpha$ systems, which are metal line systems for which $N(HI) \\ge 2 \\times 10^{20}$ cm$^{-2}$ (cf. Wolfe \\et 1986), have also received a great deal of attention in the context of what their kinematics can tell us about galaxies, particularly at high redshift. In a series of papers, Prochaska \\& Wolfe (1997,1998) present analyses of the kinematics of high redshift damped Lyman $\\alpha$ systems ($z \\simgt 2$), concluding that the ``edge leading asymmetry'' often seen in the velocity profiles of DLAs is most consistent with rotating thick disks with $v_c \\simgt 200$ \\kms. Other work has claimed that the kinematics can be explained by aggregates of small galactic fragments as would naturally be present in hierarchical structure formation models (Haehnelt, Steinmetz, \\& Rauch 1998) or simply from randomly moving clouds in spherical halos (McDonald \\& Miralda-Escud\\'e 1999), although all models appear to have problems reproducing the relative velocities of the low--ionization and high ionization species in the damped Lyman $\\alpha$ systems (Wolfe \\& Prochaska 2000). At present, little or no information on the galaxies themselves is available at these high redshifts, so that the disk inclination angles, impact parameters, and galaxy systemic redshifts are generally unknown. At lower redshifts, in contrast to the Mg~II absorbers, both ground--based and HST observations of the galaxies indicate that the galaxies associated with the high column density absorbers are a very ``mixed bag'', ranging from dwarf galaxies as faint as $0.05$ L$^{\\ast}$ to normal spirals (Steidel \\et 1994; Steidel \\et 1997; Le Brun \\et 1997; Turnshek \\et 2000; Bowen, Tripp, \\& Jenkins 2001).  In many cases spectroscopy of the putative DLA absorbers has been difficult or impossible because of the very small impact parameters (and resulting problems with scattered light from the QSOs) that often obtain for DLAs. It should be emphasized that every DLA is also a Mg~II system (cf. Rao \\& Turnshek 2000), but that the Mg~II systems are sensitive to gas with H~I column densities up to $\\sim 3$ orders of magnitude smaller than that which will produce a DLA. It is certainly possible that disk--like kinematics could dominate many Mg~II systems even when the impact parameter is well beyond the $10^{20}$ cm$^{-2}$ H~I contour. In any case, studies of systems where the absorption line kinematics and the galaxy properties can be used simultaneously to constrain kinematic models would be very valuable for understanding the DLA systems at high redshift. For the moment, this can only be done at substantially lower redshifts ($z \\simlt 1$).  The advent of echelle spectrographs on 8m--class telescopes has allowed relatively routine high dispersion spectroscopy of the same QSOs that have been used historically for Mg~II absorption surveys (e.g., Steidel \\& Sargent 1992). There has been a fair amount of activity in examining the kinematics of the intermediate redshift Mg~II systems (to scales as fine as $\\sim 5$ \\kms) and comparing them with what is known about the absorbing galaxy photometric properties and the observed impact parameters (Lanzetta \\& Bowen 1992; Churchill, Steidel, \\& Vogt 1996; Churchill \\& Vogt 2001). The relatively small samples of absorption system/absorbing galaxy pairs studied so far do not provide any obvious clues to the nature of absorbing gas in relation to the galaxies, or at least no strong systematic trends. So far, the issues of the geometry, physical conditions, and kinematics of the absorbing gas have not been considered in concert with simultaneous access to high quality imaging and spectroscopic data on the faint galaxy absorbing candidates. And yet, it is now possible with 8m-class telescopes to obtain spectra of quality sufficient to trace the kinematics of galaxy rotation curves to $z \\sim 1$ (e.g., Vogt \\et 1996), and it is clearly straightforward to obtain high quality images with kpc-scale resolution for galaxies in the same redshift range using {\\it Hubble Space Telescope} (HST).  In this paper, we present pilot observations that provide first results and explore the feasibility of establishing directly the kinematic connection between gas at large galactocentric impact parameters, and the luminous component of the galaxies.  At the very least, it should be possible to test the hypothesis that the absorbing gas is dynamically associated with the identified absorbing galaxy using greatly improved redshift and kinematic measurements of both the absorbing gas and the luminous material in the galaxy. At best, we hope to establish the nature of the absorbing gas; some possibilities include:  \\begin{itemize} \\item The gas is an extension of the galaxy disk, and the kinematics of the absorbing material are entirely compatible with a kinematic extension of the observed galaxy rotation curve. If this is the case, there may be an opportunity for measuring galaxy rotation curves out to much larger galactocentric radii than is normally possible, particularly for galaxies at relatively high redshift.  \\item The velocities of the galaxy and the absorption are too discrepant for the gas to be plausibly associated with the identified galaxy; the true absorber may remain unidentified.  \\item The galaxy and gas-phase kinematics exhibit the same general kinematic spread and zero net offset, as might be expected if the absorbing gas takes the form of distributed ``clouds'' or substructure with random velocities relative to systemic  \\item Some combination of the above. \\end{itemize}  The paper is organized as follows. In \\S 2 we present the data. \\S 3 contains a general discussion of results in each of the 5 fields, and \\S 4 explores simple kinematic models to explain the observations. Finally, in \\S 5 we discuss the general results and their implications. ",
        "conclusions": "\\subsection{Q0827+243 (OJ 248)}  The absorption system at $z_{abs}=0.52499$ is among the strongest Mg~II systems known, with the $\\lambda 2796$ component having a rest--frame equivalent width of 2.47 \\AA. The Mg~II absorption is strongly saturated, with a total velocity width of $\\simeq 270$ \\kms; some insight into the kinematic structure is possible by looking at the apparently unsaturated Mg~I $\\lambda 2853$ line, which shows at least 4 components of roughly equal strength spread over the entire velocity range covered by the Mg~II ``trough''. This system is known to be a DLA (Rao \\& Turnshek 2000), with a measured H~I column density of $2 \\times 10^{20}$ cm$^{-2}$ (which could include gas over the full $\\sim 300$ \\kms velocity range).  The spectrum of the absorbing galaxy, which was obtained with the slit oriented along the major axis of the galaxy, yields a rotation curve over the velocity range $-180 \\le v_{gal} \\le 240$ \\kms, but the blue-shifted side of the galaxy is clearly affected by a satellite galaxy that appears to be distorting the galaxy disk and which has a systemic velocity different from the galaxy of interest (skewed toward positive velocities, as can be seen in the top panel of Figure 1b). It is unclear how or if this apparent satellite is affecting the kinematics of the absorbing gas---there is no Mg~II absorption at positive velocities, but it is possible that some of the kinematic complexities of the gas may be induced by an imminent merger event. Because of this distortion to the galaxy morphology (whatever its cause), the inclination angle of the galaxy (measured to be $i=69^{\\circ}$) is rather uncertain. Assuming this inclination angle, the galaxy rotation speed is $v_c \\simeq 260$ \\kms, evaluated from the side of the rotation curve that is unaffected by the satellite (i.e., at positive velocities with respect to systemic). Unlike the other 4 galaxies discussed in this paper, G1 0827+243 has very strong [OII] emission ($W_{\\lambda}^0\\simeq 50$ \\AA) characteristic of vigorous current star formation. The rest-frame B luminosity of $\\sim 1.6$ times present--day $L^{\\ast}$, may be significantly enhanced by this star formation.  The absorbing gas kinematics are qualitatively as would be expected for a model in which the extrema of the absorbing gas velocities are consistent with a kinematic extension of the disk gas seen in emission, but extending to somewhat larger velocities at an impact parameter of $25.4$h$^{-1}$ kpc. We discuss more detailed kinematic models for this system in \\S 4.1.  \\subsection{Q1038+064 (4C 06.41)}  The absorption system at $z_{abs}=0.4415$ has been known for more than 20 years (Burbidge \\et 1977; Weymann \\et 1979) and the absorbing galaxy was among the first intermediate-redshift systems identified (Bergeron \\& Boisse\\'\\ 1991). The {\\it HST} image in Figure 2a clearly shows that the absorbing galaxy is a luminous but relatively normal mid-type spiral, with an inclination angle of $i=60^{\\circ}$. The rotation curve of the galaxy is well--determined, with an observed rotational velocity of $\\rm v_{c} sin \\theta_i \\simeq 225$ \\kms, or a de-projected rotation speed of $v_{c}^{corr} \\simeq 260$ \\kms. The absorbing gas follows the kinematics of the emitting material nicely, indicating an extreme velocity of $\\simeq 250$ \\kms relative to the galaxy systemic velocity, with the sign as expected for a simple extension of the rotation curve to a disk impact parameter of $44.8$h$^{-1}$ kpc. If the absorbing gas is interpreted as an extension of a flat rotation curve to a galactocentric radius of $\\simeq 45$h$^{-1}$ kpc, the minimum implied virial mass of the galaxy is $\\simeq 7{\\rm h}^{-1} \\times 10^{11}$ M$_{\\sun}$. We discuss the kinematic model for this system in \\S 4.2.  The Mg~I $\\lambda 2853$ absorption is very weak, with only one component securely detected at $v_{sys}=-120$ \\kms and a marginal detection of the $v_{sys}\\simeq -180$ \\kms component that is the strongest in Mg~II. Inspection of an archival HST Faint Object Spectrograph spectrum of Q1038+064 reveals a Lyman limit system at a redshift compatible with that of the Mg~II system. Assuming that the scattered light correction for the FOS spectrum is accurate, the $z=0.4415$ system has an optical depth at the Lyman limit of $\\tau_{LL} \\simeq 1.6$, or log N(H~I)$\\simeq 17.3$. Thus, there is evidence that a significant component of the extended gas giving rise to the Mg~II absorption has disk-like kinematics despite the relatively low total H~I column density.  \\subsection{Q1148+387 (4C 38.31)}  The $z_{abs}=0.5531$ absorption system was first identified by Steidel \\& Sargent 1992 (SS92). As evident in Figure 3a, the absorbing galaxy is a mid-type spiral only moderately inclined with respect to the line of sight ($\\theta_{\\rm i} \\simeq 40^{\\circ}$), and as such the amplitude of the observed rotation curve is only $\\rm v_{c} sin\\theta_i \\simeq 125$ \\kms. The (somewhat uncertain) corrected rotational velocity is $\\rm v_{c} \\simeq 195$ \\kms, consistent with or slightly low given the galaxy's luminosity. The QSO sightline, which is only 14.4h$^{-1}$ kpc in projection from the galaxy, still apparently samples only blue-shifted gas-phase velocities. However, the kinematic extent of the gas is significantly broader than the range of velocities seen from the galaxy rotation curve, and the absorbing gas appears to be distributed into $\\sim 6$ individual velocity components, all with about the same relative strength of Mg~II and Fe~II $\\lambda 2600$ absorption. We discuss kinematic models for this system in \\S 4.3.  The total H~I column density of this system can be estimated from an archival FOS spectrum. Again assuming that the zero point of the FOS flux scale is accurate, the optical depth is $\\tau_{LL} \\simeq 1.5$, or log N(H~I)$\\simeq 17.2$. This would appear to be an unexpectedly low H~I column given that the impact parameter is at a projected galactocentric distance of only $14.4h^{-1}$ kpc (corresponding to a disk impact parameter of $18.8h^{-1}$ kpc given the inclination angle). Apparently either the QSO sightline has found a ``hole'' in the H~I distribution, or the outer disk of G1 1148+387 is relatively H~I--poor.  \\subsection{Q1222+228 (Ton 1530)}  The originally-targeted Mg~II redshift along this line of sight was the $z_{abs}=0.6681$ system first discovered by Young, Sargent, \\& Boksenberg (1982). We believed, on the basis of galaxy spectra of admittedly marginal quality, that this system was produced by the nearly-edge-on spiral which was the primary target of our LRIS spectroscopy. We were somewhat surprised to find that the galaxy instead has an emission redshift of $z_{gal}=0.5502$ (leaving the much stronger $z_{abs}=0.6681$ absorption system without a confirmed absorbing galaxy candidate) from the much-higher-quality Keck/LRIS spectra. An {\\it a posteriori} search of the HIRES spectrum yielded a weak Mg~II doublet, with $W_0(\\lambda 2796) = 0.08$ \\AA, so that it is well below the detection threshold of the SS92 survey, and most earlier Mg~II absorption line surveys.  We later realized that this weak system was cataloged by Churchill \\et 1999, at $z_{abs}=0.55020$.  Ignoring for the moment the absence of an identified absorbing galaxy for the $z_{abs}=0.6681$ system (there are several candidates without spectroscopic redshifts, albeit at relatively large impact parameters, evident in Figure 4a), the $z_{abs}=0.5502$ system is interesting for a number of reasons: first, the absorption is apparently associated with a much fainter galaxy than the other 4 systems considered in this pilot study, with a circular velocity of only $v_c \\simeq 100$ \\kms. The galaxy is highly inclined ($\\theta_{\\rm i} \\simeq 75^{\\circ}$), with the projection of the major axis missing the QSO position by only 1.5 arc seconds on the plane of the sky. Despite the high inclination and the projected impact parameter of 26.5h$^{-1}$ kpc, the absorption velocity is consistent with the systemic velocity of the galaxy (see Figure 4b), and has a total velocity spread of only $\\sim 50$ \\kms.  The implications of the relative kinematics of the absorption and emission will be discussed in \\S 4.4.  There is no information on the H~I content of either the $z_{abs}=0.5502$ or the $z_{abs}=0.6681$ system, despite the existence of HST/FOS spectra, as the continuum of the QSO is cut off shortward of $\\sim 2200$ \\AA\\ by a higher redshift Lyman limit system.  \\subsection{Q1317+276 (Ton 153)}  SS92 discovered two intermediate redshift absorption systems along this line of sight, at $z_{abs}=0.2887$ and at $z_{abs}=0.6598$. An object within about 1 arc second of the QSO sightline is evident in the {\\it HST} image presented in Figure 5a, but no successful spectroscopy has been performed on this object, since the QSO is extremely bright. Earlier spectroscopy from Lick Observatory indicated that both G1 and G2 were consistent with having the same redshift as the $z_{abs}=0.6598$ absorption, within the errors (the spectra were of low quality). Because of this ambiguity, the LRIS slit was oriented so as to include both galaxies, and hence is not aligned with the major axis of either galaxy. The new spectra clearly show that G1 (which is both spectroscopically and morphologically of early type) has $z_{gal}=0.672$, which is much too high to be related to the observed absorption. Galaxy G2, on the other hand, has a systemic redshift of $z_{gal}=0.6610$, within 200 \\kms of the observed Mg~II absorption (see Figure 5b).  Galaxy G2 has a very large projected impact parameter (71.6$h^{-1}$ kpc) and because of this the identification of it as the galaxy responsible for the absorption is somewhat tentative; however, as for other more secure identifications presented in this paper, the extrema of the absorbing gas velocities are consistent with an extrapolation of the disk kinematics to the large galactocentric distances. In this case, though, attributing the gas to the disk of G2 would would imply a de-projected circular velocity of $\\simeq 250$ \\kms (which is quite reasonable for a galaxy of $\\sim 1.6L^{\\ast}$) but implies an extension of the disk--like rotation to a de-projected galactocentric radius of $\\sim 130$h$^{-1}$ kpc, possibly stretching the bounds of feasibility\\footnote{The circular velocity implied for G2 is also consistent with the measured rotation curve, although the fact that the slit was placed at an angle of $\\sim 60$ degrees relative to the major axis makes the de-projection of the rotation curve somewhat model-dependent.}.  The absorbing gas yields a clear detection of Mg~I associated with the dominant component of Mg~II at $v_{sys}=-170$ \\kms, and a hint of Mg~I in the weaker $v_{sys}=-80$ \\kms component. There does appear to be significant Mg~II absorption all the way to $v_{sys}=0$. This is circumstantial evidence that the gas is indeed dynamically related to galaxy G2. We will discuss this system in more detail in \\S 4.5.  There is a strong Lyman $\\alpha$ absorption line ($W_{\\lambda}^0=1.5$ \\AA) associated with the Mg~II absorption, and an optically thick Lyman limit (measured by Bahcall \\et 1993 at $\\tau = 3.5$ and Churchill \\et 2000 at $\\tau = 5.4$. Given uncertainties in the flux zero point of FOS data, we take these measurements as lower limits, suggesting an H~I column density of $\\simgt 10^{18}$ cm$^{-2}$). Interestingly, there is a strong complex of Lyman $\\alpha$ forest lines extending from $z=0.660$ to $z=0.672$ ($\\sim 2200$ \\kms; Bahcall \\et 1996) and there are at least 2 galaxies (including the early type G1 1317+276) near the line of sight at $z \\simeq 0.67$.    We have presented data on 5 intermediate redshift Mg~II absorption systems for which there is much more information than has been available in the past, and in many ways it has made for a more puzzling picture of the nature of absorbing material in the outer parts of (in these cases, disk) galaxies. Most recent models for the kinematics of Mg~II absorption systems have involved two separate components contributing to the kinematics-- a rotating disk component, which is thought to produce the ``dominant'' components of complex kinematic systems (e.g. Charlton and Churchill 1998, Churchill \\& Vogt 2001), with more symmetrical (with respect to the systemic redshift) ``halo'' components providing a broad distribution of velocity, whose origins might be ascribed to random motions, infall, or outflow. Given the limited geometrical information from Mg~II--selected galaxy surveys (Steidel \\et 1997, SDP, Bergeron \\& Boiss\\'e\\ 1991) at these redshifts, and the limited morphological information on the absorbing galaxies (Steidel 1998; Steidel \\et 1997), this picture has the advantage of being compatible with both the spiral nature of most of the galaxies, and the large (and possibly quasi-spherical) gaseous envelopes surrounding a very large fraction of $z \\sim 0.6$ field galaxies within $\\sim 1.5$ magnitudes of present-day $L^{\\ast}$.  However, based on the systems presented in this paper, the situation cannot be so simple-- the components that are not easily explained by thin disk rotation must have the systematics that are like those produced by rotation. There is clear evidence for rotation in 4 out of the 5 cases, since not only is all of the absorption offset to one side of the systemic redshift, in all cases it is in the right sense to be qualitatively explained by an extension of the disk rotation to the line of sight (which is well beyond where the optical rotation curves are measured, by factor of 2 to as much as a factor of 6). As detailed in \\S 4, though, given knowledge of the disk inclinations and impact parameters (not known for previous analyses of Mg~II kinematics), disk--like rotation is not enough to explain the bulk of kinematic components seen in absorption.  Our simple kinematic modeling of rotating ``thick disks'' in \\S 4 (which are by no means intended to be unique solutions to the kinematic conundrum) imply that for two of the galaxies (G1 1148+3842 and G2 1317+276) a successful model does not involve an extension of the disk at all, but requires essentially a {\\it rotating halo}. In the three other cases, the observed kinematics are reasonably well fit by thick disk models in which the circular velocities are a rapidly decreasing function of scale height, with the extrema of the velocities being produced at the disk intersection, but where the gas at lower relative velocities with respect to systemic comes from fairly large $z$ distances. As discussed in \\S 4, the galaxies best--fit with the velocity scale height models allow for a wide range of solutions so long as the ratio $H_{\\rm eff}/h_v$ remains roughly constant, as there are only weak constraints on the required thickness of the gas layer above the extrapolation of the plane of the disk. Given what is known about the geometry of the gaseous envelopes capable of giving rise to easily detectable Mg~II absorption, we would be inclined to favor models in which $H_{\\rm eff}$ is on the order of $\\sim 30-40h^{-1}$ kpc for $L^{*}$ galaxies like G1 1038+064 and G1 0827+243. On the other hand, G2 1222+228, which is among the faintest Mg~II absorbers identified at comparable redshift, must be quite different from these larger spirals, in that the effective thickness of the disk must be quite thin to {\\it avoid} producing absorption at large velocities with respect to systemic, and the velocity scale height must be even smaller to bring the kinematic model into tolerable agreement with the data.  It is instructive to consider possible local analogs of the kinematic behavior of galactic gas we observe at $z \\sim 0.5$. For highly ionized gas in the Galaxy seen in absorption against the continua of hot stars and extragalactic AGN, a model in which the ``halo'' gas co-rotates with the disk up to heights of several kpc above the plane has often been assumed, and recent observations seem to require this (Savage, Sembach, \\& Lu 1997).  The scale heights reached appear to vary as a function of the ionization state of the ion, with the most highly-ionized species extending to the largest distances above the plane. However, in at least 2 of the cases we have observed, Mg~II would have to have an effective scale height that is more than an order of magnitude larger than that of C~IV in the Galaxy, and in the other cases a completely co-rotating halo would fail badly to reproduce the observed absorption line kinematics.  A recent development based on the most sensitive H~I (e.g., Swaters, Sancisi, \\& van der Hulst 1997; Sancisi \\et 2000, Schaap, Sancisi, \\& Swaters 2000) and H~II (e.g., Rand 2000) measurements locally is the ability to follow the kinematics of the gas to relatively large $z$ distances above the planes of spiral galaxies. Rand (2000) finds that for the edge-on starburst galaxy NGC 5775, there is gas whose rotation velocity has decreased to zero by a height of 5 kpc. He interprets this behavior as a trend of decreasing rotational velocity as a function of $z$, similar to our modeling above. Swaters \\et (1997) observed clear evidence for a systematically smaller rotation velocity of the ``H~I halo'' of NGC 891, with gas at several kpc above the plane rotating 25--50 \\kms more slowly than in the plane. Sancisi \\et (2000) discuss 21-cm observations of galaxies with ``beards'', in which gas observed away from the plane of the galaxy seems to ``know'' about the rotation of the disk (i.e., the rotation has the same general direction as far as can be discerned with the observations) but has kinematics that represent large departures from the disk rotation. The qualitative similarities of these observations to those presented in this paper are clear; however, it is unclear how common the kinematically ``anomalous'' gas is in local spirals (only a few galaxies have been observed to the required level of sensitivity). In any case, the gas-phase kinematics of $z \\sim 0.5$ galaxies refer to much larger galactocentric distances and much smaller H~I column densities (with the exception of G1 0827+243 which would easily be observed in 21 cm emission if it were nearby).  The interpretation of the slowly rotating gas in nearby galaxies is largely qualitative at the present time-- the authors cited above invoke both hydrodynamic disk/halo cycling of gas, and changes in the gravitational potential with $z$ distance above the plane, as possible explanations of the observations. If the higher redshift objects are at all analogous, it is hard to imagine that the gravitational potential argument can be relevant, since the sight lines all intersect the galaxies at radii where dark matter would be expected to dominate strongly over a baryonic disk.  Bregman (1980) considered kinematic models of the disk/halo circulation (the ``Galactic fountain'' model) in which parcels of hot gas are expelled from the disk by star formation events and may, by the time they cool, have been transported to both a large height above the plane and to a larger distance from the galaxy rotation axis. Because of conservation of angular momentum, the gas at large heights above the galactic plane would lag with respect to the disk rotation, producing a more slowly-rotating ``halo''. In practice, the kinematics of the gas as a function of height $z$ would depend on the details of the distribution and energetics of previous star-formation episodes and on the pressure profile of the galaxy as a function of galactocentric distance; in principle, gas with any velocity between systemic and $v_{c} sin\\theta_{\\rm i}$ could be observed at any position along the line of sight. Nevertheless, as long as the rotation of the gas dominates over radial motions (infall or outflow) such a picture would be qualitatively consistent with our observations.  In the nearby starburst galaxies discussed above, very active current star formation lends qualitative credence to fountain flows as a possible explanation of the observed gas-phase kinematics, but this leads to a puzzle for the higher redshift objects-- with the exception of G1 0827+243, none of the galaxies considered here has an unusually high rate of star formation, and certainly there is no active star formation coinciding with the large inferred disk impact parameters. The observation that the Mg~II absorbers at intermediate redshift tend in general {\\it not} to be particularly active star-formers has been used previously to argue against fountain-type flows being important to the presence of extended gaseous envelopes (SDP, Steidel 1995), but this argument assumes that the timescale for the circulation of the gas is relatively short. We now consider the possibility that past star formation, now observed only through the older stellar populations present in the galaxies, might be responsible for the rotating halo gas observed at $z \\sim 0.5$.   There is mounting evidence for the importance of large-scale galactic winds for star forming galaxies at high redshift. Observations of $z \\simgt 3$ Lyman break galaxies show clearly that the strong far-UV interstellar absorption lines are blue-shifted, and the Lyman $\\alpha$ emission lines red-shifted, by up to 1000 \\kms with respect to systemic (Pettini \\et 1998, 2001), with typical implied outflow velocities being several hundred \\kms. Very recently, it has been shown that the properties of the Lyman $\\alpha$ forest are strongly affected by the presence of LBGs $z \\sim 3$, and that several different observations can be explained simultaneously if the super-winds have a sphere of influence of $\\sim 125h^{-1}$ kpc on average (Adelberger \\et 2001). The cooling time for the shock heated gas in the halo can be very long, and it is at least conceivable that this gas, or similar gas ejected at lower redshifts (where there are currently fewer observations constraining the extent of super-winds), could ultimately supply the gas that forms the bulk of the disk observed at $z \\sim 0.5$ and the material that produces Mg~II absorption. The physics of the multi-phase gas that no doubt results from the wind activity is very complex, and a full treatment is well beyond the scope of this paper. It is not clear that this kind of flow would really result in a ``memory'' of disk-like kinematics in the halo, since the material involved in the type of super-winds inferred to exist at $z \\sim 3$ would tend to originate from low angular momentum gas that has settled to a very compact nuclear region where most of the star formation appears to take place. An alternative possibility is that the gas has acquired significant angular momentum during the extended time that it spends at large galactocentric distances, and that this angular momentum is naturally strongly correlated with that of the gas that has found its way to the disk. In this scenario, much of the gas falling onto the disk would do so gradually, would be significantly metal-enriched, and would have the same kinematic systematics as disk gas, albeit with smaller rotational velocities. It is not entirely clear whether the observed kinematics at $z \\sim 0.5$ are consistent with both rotation {\\it and} infall; the  sample is too small to justify a more in-depth treatment at this time. In any case, a general picture of halo gas being ``recycled'' disk gas may actually help explain the loose correlation of inferred size of the Mg~II-absorbing envelope with stellar mass, the inferred roughly axisymmetric geometry of the envelope, and the persistence of absorbing gas over much longer than the typical galaxy dynamical time of $\\sim$ a few $\\times 10^8$ years.    It is almost certainly premature to generalize about the nature of the Mg~II absorbing gas, given the small sample of 5 systems presented here and the fact that each system requires somewhat different assumptions to find kinematic models that are adequate. We have not considered in detail whether any of these ad hoc kinematic models are physically plausible, and we have not considered the hydrodynamics of the gas at all. More detailed modeling seems unjustified until a larger sample is in hand.  It is clearly worth extending this type of study in two ways. First, it is essential to obtain accurate redshifts and (where possible) rotation curves for a larger sample of Mg~II--selected galaxies (the HIRES and WFPC-2 data are already on hand for a sample of $\\sim 25$ Mg~II absorbers). Secondly, it will be important to obtain high quality absorption line data extending to higher ionization species (like C~IV) to see if the rotation signatures are as clear in the highly ionized component as they are for the gas presently traced by Mg~II. Initial forays in this direction have already been made by Churchill \\et (2000), but the archival FOS data are generally of too-coarse resolution to compare absorption line kinematics in detail. On the other hand, a fraction of the QSOs in the sample for which we have WFPC-2 images are bright enough for STIS high dispersion spectroscopy in reasonable integration times, and this should prove a fruitful line of research in the future.  \\bigskip \\bigskip We would like to thank Kurt Adelberger for help with the observations and for many discussions. Useful conversations with Betsy Barton-Gillespie, Liese van Zee, and Jason Prochaska are gratefully acknowledged. CCS and AES have been supported in part by grants AST95-96229 and AST-0070773 from the U.S. National Science Foundation and by the David and Lucile Packard Foundation. Early work on the HST data was supported by grant GO-05984.01-94A and GO-06577.01-95A from the Space Telescope Science Institute."
    },
    "0201/astro-ph0201165_arXiv.txt": {
        "abstract": "The NGC~4410 group of galaxies provides us a rare opportunity to study a nearby (97 $h_{75}^{-1}$ Mpc) example of a radio galaxy (NGC~4410A) embedded in an extended X-ray source, with evidence for star formation that can be readily spatially distinguished from regions dominated by the AGN and shocks. We present broadband and narrowband optical images along with optical and IUE ultraviolet spectroscopy for the radio galaxy NGC 4410A and its companion NGC 4410B.   Our H$\\alpha$+[N~II] images reveal six luminous H~II regions (L$_{H\\alpha}$ $\\sim$ 10$^{40}$ erg s$^{-1}$) distributed in an arc near NGC 4410A. Partially completing the ring is a prominent stellar loop containing diffuse ionized gas.  This filamentary gas, in contrast to the H~II regions, shows spectroscopic signatures of shock ionization. The star formation in this system may have been triggered by a collision or interaction between the two galaxies, perhaps by an expanding density wave, as in classical models of ring galaxies. Alternatively, the star formation may have been induced by the impact of a radio jet on the interstellar matter. Extended Ly$\\alpha$ is detected in the ultraviolet IUE spectrum. The ultraviolet continuum, which is presumably radiated by the nucleus of NGC~4410A, is not extended. NGC~4410A appears to be interacting with its neighbors in the NGC~4410 group, and could be an example of a spiral galaxy transforming into an elliptical. ",
        "introduction": "}  We do not understand the link between the radio sources in active galactic nuclei and the emission line regions surrounding those sources. The extended emission line regions around powerful radio galaxies may be radiating light scattered from a central anisotropic source (e.g. Tadhunter et al. 1993; Fabian 1989). Alternatively this gas could be heated by shocks produced by the radio jet interacting with the interstellar medium (ISM), or some combination of shock heating and star formation (Rees 1989; Begelman \\& Cioffi 1989; De Young 1989). In the centers of cooling flow clusters, the radio galaxy morphology seems to be intimately related to the structures seen radiating optical and near-infrared emission lines typical of star-forming regions or gentle shocks (Donahue et al. 2000; Koekemoer et al. 2000). In the cooling flow clusters, the radio source seems to either excavate a cavity in the emission-line gas or the emission-line gas forms around the radio cavity. Similar structures are seen at high resolution in the latest Chandra images of the central elliptical galaxies in cluster cooling flows such as Hydra A (McNamara et al. 2000). Given the variety of environments of these emission line regions, more than one process may be responsible for the emission-line nebulae in radio galaxies and central galaxies in clusters.  To investigate these issues, we made detailed studies of a nearby radio galaxy in a group of galaxies that is close enough to resolve individual H~II regions, the peculiar low-luminosity radio galaxy NGC 4410A. This system is quite nearby (97 $h_{75}^{-1}$ Mpc for $H_0 = 75h_{75}$ km s$^{-1}$ Mpc$^{-1}$), and so provides a good target for detailed studies. NGC 4410A, which is classified as an Sab? Pec galaxy on the basis of its optical appearance (de Vaucouleurs et al. 1991), has a prominent bulge surrounded by an extended ring or loop (Hummel, Kotanyi, $\\&$ van Gorkom 1986).  It forms a close pair with the nearby S0? Pec galaxy NGC 4410B. These two galaxies are part of a sparse group containing a dozen known members (Smith 2000). The velocity dispersion of those members is $\\sim220\\pm70$ km s$^{-1}$, from the velocities reported in Smith (2000), and the group has no obvious group emission (Tsch\\\"oke et al. 1999). Therefore this group is a poor, X-ray faint group (e.g. Zabludoff \\& Mulchaey 1998), a classification that is consistent with its lack of elliptical galaxies. The radio luminosity of NGC 4410A is near the faint end of the luminosity range for radio galaxies (Condon, Frayer, \\& Broderick 1991), with a 4.8 GHz luminosity of 1.5 $\\times$ 10$^{23} h_{75}^{-2}$ W Hz$^{-1}$. The far-infrared (42.5 $-$ 122.5 $\\mu$m) luminosity of the NGC 4410A+B pair is also moderate, 3.9 $\\times$ 10$^9 h_{75}^{-2} \\rm{L}\\sun$ (Mazzarella, Bothun, \\& Boroson 1991, calculated as in Lonsdale et al. 1985). The $L_{FIR}/L_{4.8~\\rm{GHz}}$ ratio for this system is about 200 times less than expected for a galaxy dominated by star formation, confirming that NGC 4410A contains a radio-loud active nucleus (Condon et al. 1991), in spite of its optical classification as a disk galaxy.  NGC 4410A has a very peculiar radio morphology: a shorter (3$'$ $\\sim 80 h_{75}^{-1}$ kpc) very distorted radio lobe to the southeast (Hummel et al. 1986; Batuski et al. 1992) and a fainter longer structure (7$'$ $\\sim 200 h_{75}^{-1}$ kpc) to the northwest (Smith 2000). This strange radio morphology may have been caused by an interaction of the radio lobe with interstellar matter that has been disturbed by a gravitational encounter between NGC 4410A and NGC 4410B (Smith 2000). NGC 4410A+B contains abundant interstellar matter ($M_{HI} \\sim 10^9 h_{75}^{-2} \\msun$ and $M_{H_2} \\sim 4 \\times 10^9h_{75}^{-2} \\msun$ (Smith 2000; assuming the standard Galactic $I_{CO}/N_{H_2}$ conversion factor.) About a third of the HI lies in a tail-like structure extending 1\\farcm7 ($50 h_{75}^{-1}$ kpc) to the southwest, coincident with a faint optical tail (Smith 2000). This tail overlaps with the southeastern radio lobe, suggesting an interaction between the radio jet and the HI gas.  In addition to an active nucleus, NGC 4410A  has on-going star formation. In this paper, we present evidence for luminous H~II regions in NGC 4410A, in the form of new optical images and optical and ultraviolet spectroscopy. ",
        "conclusions": "NGC 4410A is the host of both a radio-loud active nucleus and on-going star formation. The star formation is in the form of extremely luminous H~II regions aligned in an arc along the eastern portion of a ring-like stellar structure surrounding the active nucleus. The western portion of this ring contains filamentary ionized gas with optical line ratios implying shock ionization. NGC~4410A is thus a nearby example of an AGN with spatially-resolved evidence for both star formation and shocks in an extended emission line region. Such an object lends credence to the hypothesis that the emission line regions associated with radio galaxies and the central galaxies in cluster cooling flows might also result from the effects of more than one phenomenon, with the implication that no single physical process can explain all of the observations. NGC 4410A could be an example of a disk galaxy on its way to becoming an elliptical, via interactions with its group companions."
    },
    "0201/astro-ph0201486_arXiv.txt": {
        "abstract": "{Here we attempt to infer the recent history of star formation in the BCD galaxy VII Zw403, based on an analysis that accounts for the dynamics of the remnant generated either by an instantaneous burst or by a continuous star formation event. The models are restricted by the size of the diffuse X-ray emitting region, the H$_{\\alpha}$ luminosity from the star-forming region and the superbubble diffuse X-ray luminosity.  We have re-observed VII Zw403 with a better sensitivity corresponding to the threshold H$_{\\alpha}$ flux $8.15 \\times 10^{-17}$ erg cm$^{-2}$ s$^{-1}$. The total H$_{\\alpha}$ luminosity derived from our data is much larger than reported before, and presents a variety of ionized filaments and incomplete shells superimposed on the diffuse H$_{\\alpha}$ emission. This result has a profound impact on the predicted properties of the starburst blown superbubble. Numerical calculations based on the HST H$_{\\alpha}$ data, predict two different scenarios of star formation able to match simultaneously all observed parameters. These are an instantaneous burst of star formation with a total mass of $5 \\times 10^5$ M$_{\\odot}$ and a star-forming event with a constant SFR = $4 \\times 10^{-3}$ M$_{\\odot}$ yr$^{-1}$, which lasts for 35 Myrs. The numerical calculations based on the energy input rate derived from our observations predict a short episode of star formation lasting less than 10 Myrs with a total star cluster mass $\\sim (1 - 3) \\times 10^6$ M$_{\\odot}$. However, the five main star-forming knots are sufficiently distant to form a coherent shell in a short time scale, and still keep their energies blocked within local, spatially separated bubbles. The X-ray luminosities of these is here shown to be consistent with the ROSAT PSPC diffuse X-ray emission.} ",
        "introduction": "It has recently been recognized that the star formation activity in galaxies is very irregular in time, and many examples of major burst episodes exhibit an extremely high star formation rate concentrated in well localized space regions (Terlevich 1996). It is also now well known that starbursts (SBs) cause an emission that dominates the entire host galaxy luminosity and their mechanical energy input rate is expected to cause major structural changes in the surrounding  interstellar medium (ISM). In this respect it has become of great interest to study the properties of the resultant large-scale expanding superbubbles which, powered by the violently injected newly processed matter, establish the time scale for mixing with the ISM (Tenorio-Tagle 1996, Silich et al. 2001). In extreme cases, the superbubbles are thought to break out of the galactic discs leading to an effective mass and energy transport into the low density halos or even into the intergalactic medium via a superwind (Heckman \\etal 1990).  Starbursts in the local universe are also assumed to be good representatives of the star-forming activity at high redshifts. This concept defines their cosmological interest as key laboratories for studying the ISM, the transport of supernovae processed metals, as well as the chemical evolution of galaxies and of the intergalactic medium. The resulting structure in the ISM due to mechanical energy injected by SBs is very similar to the interstellar wind-blown bubbles around single massive stars (see Weaver \\etal 1977 for their four zone model), although the much larger energy input rate in SBs leads rapidly to much larger scales. Hydrodynamical simulations (see Tenorio-Tagle \\& Bodenheimer 1988, Bisnovatyi-Kogan \\& Silich 1995 and references therein; Suchkov \\etal 1994; Silich \\& Tenorio-Tagle 1998; D'Ercole \\& Brighenti 1999; Strickland \\& Stevens 2000) currently include differential galactic rotation, radiative cooling, strong density gradients between the disk and the halo and thus are able to follow the moment of breakout, as well as the fragmentation of the expanding outer shell via Rayleigh - Taylor instabilities and the venting of the superbubble hot interior gas, either into the intergalactic space, or into the host galaxy halo (the blowout phenomenon).  Most of the up to-date simulations have been performed under the assumption of a constant energy deposition rate, as expected from an instantaneous burst model. However, studies of the stellar population in OB associations related to young ($\\tau_{OB} < 10$ Myr) Large Magellanic Cloud (LMC) bubbles (see  Oey \\& Smedley, 1998 and references therein) have demonstrated that ``realistic'' energy input rates are very different from the assumed constant energy input rates used in numerical simulations. Thus the instantaneous burst assumptions may not be applicable to all cases. Here we attempt to establish a method of comparison between the theory of superbubbles and the observations of remnants produced by massive star formation in galaxies. Two possible modes of star formation, instantaneous and extended bursts, are taken into consideration. For both cases the mechanical luminosity, ultraviolet photon output, mass returned to the ISM, and the fraction of each in metals, all as a function of time,  are estimated. On the other hand we have the observed parameters: the H$_{\\alpha}$ or H$_{\\beta}$ luminosity which can be directly related to the SFR under the assumption that all photons are used up in the ionized region. One can also estimate the size and luminosity of the X-ray remnant. In some cases the remnants may have slowed down sufficiently to display their outer expanding shells, either in the optical or in HI observations, or both, giving further information about the size, expansion speed, and mass behind the outer shock. A comparison with the theory also requires some preconception of the galaxy's ISM, which can as a first approximation be derived from HI observations and the inferred dynamical mass. However, one also needs to make an assumption about the fraction of this gas locked up in dense clouds and immersed into a less dense  ISM background.  With the aim of establishing a method to confront theory with observations that may lead us to infer the recent star formation history of galaxies, the blue compact dwarf (BCD) VII Zw403 is thoroughly analyzed. Section 2 derives the main properties of coeval and extended star formation modes. Section 3 summarizes the main observational properties of our target. Section 4 presents a summary of the hydrodynamical calculations aimed at matching the observed properties. The results of the calculations and our main findings are discussed in Section 5. ",
        "conclusions": "Here we have discussed the recent history of star formation and a possible nature of the diffuse X-ray emission in the nearby BCD galaxy VII Zw403. Two possible scenarios of star formation have been considered: an instantaneous burst, and an extended episode of star formation.  To construct the numerical model we have provided new narrow-band observations of VII Zw403 centered on the H$_{\\alpha}$ line with a long exposure time corresponding to the threshold H$_{\\alpha}$ flux $8.15 \\times 10^{-17}$ erg cm$^{-2}$ s$^{-1}$. These observations reveal a variety of ionized filaments and incomplete shells superimposed on the diffuse H$_{\\alpha}$ emission that most certainly result from the photons leaking out of the main star-forming centers. The largest feature is the 250 pc broken shell associated with stellar association 4. The total H$_{\\alpha}$ luminosity derived from our observations, L$_{H_{\\alpha}} = 4.7 \\times 10^{40}$ erg s$^{-1}$, is much larger than reported before. This has a profound impact on the predicted properties of the starburst blown superbubble.  The numerical models based on the HST H$_{\\alpha}$ data require either an instantaneous burst of star formation with a total mass of $5 \\times 10^5$ M$_{\\odot}$, or a star formation episode with a constant SFR = $4 \\times 10^{-3}$ M$_{\\odot}$ yr$^{-1}$ lasting 35 Myr. The models however require radically different structures of the galactic ISM and imply very different properties  of the resulting remnant.  The best coeval model assumes most of the ISM to be locked up within high density clouds, and only $\\sim 5\\%$ of the observed neutral hydrogen mass is in the smooth component. The hydrodynamical calculations also predict the outer shell to be adiabatic after reaching a 1 kpc radius, and to contribute 50 - 80\\% of the observed diffuse X-ray emission. The bubble evolutionary time is estimated to be $\\tau_{dyn} \\approx 7$ Myr when its expansion speed is $\\approx 200$ km s$^{-1}$.  The best continuous star formation model requires of much higher density in the smooth ISM component, with only $\\sim 30\\%$ of the HI mass concentrated in dense clouds. This leads to a much smaller bubble expansion velocity (V$_{exp} \\approx 35$ km s$^{-1}$), larger evolutionary time ($\\tau_{dyn} \\sim 35$ Myr), and a rapid cooling within the outer shell. That is, this model predicts a low brightness HI shell surrounding the diffuse 1 kpc X-ray region. This shell may also show up in ${H\\alpha}$ if exposed to the $UV$ flux from the central cluster. The inner gas metalicities are also predicted to be very different in these two cases.  The numerical calculations based on the high energy input rate derived from our observations require an instantaneous burst or a short episode of star formation with SFR $\\sim 0.1$ M$_{\\odot}$ yr$^{-1}$ lasting less than 10 Myr with similar total stellar cluster masses $(1 - 3) \\times 10^6$ M$_{\\odot}$ and most of the ISM ($\\sim 95\\%$) locked up within high density clouds. The comparison of the energy input rate derived from our ${H\\alpha}$ data with the theoretical limits, implies that the entire ISM and metals produced by the current episode of star formation are going to be ejected from the galaxy after the coherent superbubble is formed.  It appears that the five main star-forming knots are sufficiently distant to form a coherent shell in a short time scale, while keeping their energies blocked within local, spatially separated bubbles. This provides a time delay that must be considered when developing a numerical model for the coherent superbubble driven by a number of young stellar clusters. Numerical calculations show that the X-ray luminosities from young local bubbles are in a good agreement with the ROSAT PSPC data. This agreement indicates that the observed diffuse component of the X-ray emission may be related to the small centrally concentrated bubbles, rather than to the coherent 1 kpc structure.  Further observations with the XMM-NEWTON observatory is expected to be able to recover the real nature of the diffuse X-ray emission and the recent history of star formation in this galaxy.  We thank D. Bomans for his comments and suggestions regarding the X-ray data. We also thank the anonymous referee for his detailed report that greatly improved our paper. Finally we also thank Edward Chapin for his careful reading of the manuscript.  This work has been supported by the Spanish grants PB97-1107 and AYA2001-3939, and the Mexico (CONACYT) project 36132-E."
    },
    "0201/astro-ph0201435_arXiv.txt": {
        "abstract": "We estimate the distribution of apparent axis ratios $q$ for galaxies in the Sloan Digital Sky Survey (SDSS) Early Data Release. We divide the galaxies by profile type (de Vaucouleurs versus exponential) as well as by color ($u^* - r^* \\leq 2.22$ versus $u^* - r^* > 2.22$). The axis ratios found by fitting models to the surface photometry are generally smaller than those found by taking the second moments of the surface brightness distribution. Using the axis ratios found from fitting models, we find that galaxies with de Vaucouleurs profiles have axis ratio distributions which are inconsistent, at the 99\\% confidence level, with their being a population of randomly oriented oblate spheroids. Red de Vaucouleurs galaxies are slightly rounder, on average, than blue de Vaucouleurs galaxies. By contrast, blue galaxies with exponential profiles appear very much flatter, on average, than red galaxies with exponential profiles. The red exponential galaxies are primarily disk galaxies seen nearly edge-on, with reddening due to the presence of dust, rather than to an intrinsically red stellar population. ",
        "introduction": "Observationally based estimates of the three-dimensional shapes of galaxies serve as a diagnostic of the physics of galaxy formation and evolution. Galaxies can only be seen in projection against the sky; thus, astronomers can only attempt to deduce their three-dimensional properties from their two-dimensional, projected properties. Obviously, information about intrinsic shapes is lost in projection; for instance, using only the two-dimensional surface photometry of a given galaxy, it is impossible to determine its intrinsic three-dimensional shape.  Elliptical galaxies have isophotes that are well approximated as ellipses (hence the name ``elliptical''). The shape of an ellipse is specified by its axis ratio $q$, with $0 \\leq q \\leq 1$. The three-dimensional isophotal surfaces of elliptical galaxies are generally modeled as ellipsoids. A stellar system whose isophotal surfaces are similar, concentric ellipsoids, without axis twisting, will have projected isophotes which are similar, concentric ellipses, without axis twisting \\citep{co56,st77}. The apparent axis ratio $q$ of the projected ellipses depends on the viewing angle and on the intrinsic axis ratios $\\beta$ and $\\gamma$ of the ellipsoid. Here, $\\beta$ is the ratio of the intermediate to long axis, and $\\gamma$ is the ratio of the short to long axis; thus, $0 \\leq \\gamma \\leq \\beta \\leq 1$.  Beginning with \\citet{hu26}, many attempts have been made to deduce the distribution of intrinsic shapes of elliptical galaxies, given their distribution of apparent shapes. The early assumption, in the absence of evidence to the contrary, was that elliptical galaxies were oblate spheroids, flattened by rotation \\citep{sa70}. If elliptical galaxies were all oblate spheroids ($\\beta = 1$), or all prolate spheroids ($\\beta = \\gamma$), and if their orientations were random, then it would be possible to deconvolve their distribution of apparent axis ratios $f(q)$, to find their distribution of intrinsic axis ratios $N(\\gamma)$. However, the pioneering work of \\citet{be75} and \\citet{il77} led astronomers to abandon the assumption that elliptical galaxies are necessarily oblate. The shapes of ellipticals have been reanalyzed with the assumption that they are intrinsically prolate or triaxial, rather than oblate \\citep{bi78, be80, bi80, bi81, ry96}.  Statements about the intrinsic shapes of galaxies must be statistical in nature, since astronomers do not exactly know the distribution $f(q)$ of axis ratios for a given class of galaxy. In this paper, we will be examining the apparent axis ratios for galaxies in the Sloan Digital Sky Survey(hereafter SDSS; York {\\em et al.} 2000). The set of axis ratios we analyze constitute a finite sample drawn from a parent population $f(q)$. In this paper, we take into account the finite size of the sample in rejecting or accepting, at a known confidence level, two null hypotheses: that the galaxies are randomly oriented oblate spheroids, or that they are randomly oriented prolate spheroids. To accomplish this, we make a kernel estimate, ${\\hat f}(q)$, of the distribution of axis ratios and mathematically invert ${\\hat f}(q)$ to find ${\\hat N_o}(\\gamma)$ and ${\\hat N_p}(\\gamma)$, the estimated distribution of intrinsic axis ratios for a population of oblate spheroids and a population of prolate spheroids, respectively.  The rest of the paper is organized as follows. In section~\\ref{data}, we describe the Sloan Digital Sky Survey, and the methods by which the apparent axis ratios of the galaxies are estimated. In section~\\ref{method}, we present a brief review of the nonparametric kernel estimators used in this paper. In sections~\\ref{deV} and \\ref{exp}, we find the kernel estimate ${\\hat f}(q)$ for galaxies with de Vaucouleurs luminosity profiles and for galaxies with exponential profiles, and find the implications for their intrinsic shapes. In section~\\ref{discuss}, we discuss our results. ",
        "conclusions": "\\label{discuss} The galaxies in the SDSS EDR with de Vaucouleurs profiles have a distribution of apparent shapes which is incompatible (at the 99\\% confidence level) with their being randomly oriented oblate spheroids. This is consistent with the result found by \\citet{la92} for a sample of 2135 elliptical galaxies with shapes estimated from survey plates of the APM Bright Galaxy Survey. When the SDSS survey is complete, it will provide a sample of galaxies $\\sim 20$ times larger than the SDSS EDR. This increase in sample size will enable us to determine more accurately the distribution of apparent axis ratios $f(q)$. The kernel width $h$ will be decreased by a factor $\\sim 20^{-0.2} \\sim 0.55$. The error intervals, which are essentially determined by the $N^{1/2}$ fluctuations in bins of width $h$, will be reduced by a factor $\\sim 20^{-0.4} \\sim 0.30$. Although a simple increase in the sample size will not enable us to determine the true distribution of intrinsic shapes, it will enable us to make stronger statistical statements about our rejection or acceptance of the prolate or oblate hypothesis.  The blue de Vaucouleurs galaxies, with a mean axis ratio of $\\langle q_{\\rm model} \\rangle = 0.639$, are only slightly flatter in shape than the red de Vaucouleurs galaxies, with $\\langle q_{\\rm model} \\rangle = 0.652$. Thus, if the color of the blue de Vaucouleurs galaxies is the result, at least in part, of recent star formation, we can conclude that the overall shape of the galaxies is not strongly affected by star formation. Although elliptical galaxies with old stellar populations tend to be rounder than those with young stellar populations \\citep{ry01}, this difference is only large at small radii ($r \\lesssim r_e/8$), while the values of $q_{\\rm model}$ used in this paper emphasize the axis ratio at much larger radii ($r \\gtrsim r_e$).  The galaxies with exponential profiles have, by contrast, shapes which are strongly dependent on color, with the red exponential galaxies consisting predominantly of dust-reddened edge-on (or nearly edge-on) disks. A significant number of galaxies in the SDSS EDR appear to be edge-on late-type galaxies, with exponential profiles, rather than early-type galaxies, with de Vaucouleurs profiles. The number of red exponential galaxies ($N = 815$) is 14.4\\% of the number of red de Vaucouleurs galaxies ($N = 5{,}659$). Thus, if we attempted to select out elliptical galaxies purely on the basis of color, we would have been faced with a significant contamination by reddened disks. Spectroscopy or accurate surface photometry are required to distinguish between elliptical galaxies and disk galaxies. This analysis agrees very well with the result of \\citet{sc99} and carries substantially more statistical weight.  In summary, galaxies with de Vaucouleurs profile have an axis ratio distribution consistent, at a high confidence level, with their being randomly oriented prolate spheroids (though it is also consistent with their being triaxial systems). Galaxies with exponential profiles have an axis ratio distribution which is dependent on color, suggesting that red exponential galaxies are nearly edge-on systems reddened by dust. Since a fair number of red galaxies in the SDSS EDR are nearly edge-on exponential disks, it is dangerous to select elliptical galaxies purely on the basis of color."
    },
    "0201/astro-ph0201059_arXiv.txt": {
        "abstract": "{ {\\em Chandra} has resolved the starburst nuclear region of the face-on grand-design spiral M83. Eighty-one point sources are detected (above 3.5-$\\sigma$) in the ACIS S3 image, and 15 of them are within the inner 16\\arcsec~region of the galaxy. A point source with $L_{\\rm x} \\approx 3 \\times 10^{38}$~erg~s$^{-1}$ in the 0.3--8.0 keV band is found to coincide with the infra-red nuclear photometric peak, one of the two dynamical nuclei of the galaxy. No point-like sources are resolved (at a 2.5-$\\sigma$ level) at the centre of symmetry of the outer optical isophote ellipses, suspected to be another dynamical nucleus. About 50\\% of the total emission in the nuclear region is unresolved; of this, about 70\\% can be attributed to hot thermal plasma, and the rest is probably due to unresolved point sources (eg, faint X-ray binaries). The azimuthally-averaged radial distribution of the unresolved emission has a King-like profile, with no central cusp. Strong emission lines are seen in the spectrum of the optically thin plasma component. The high abundances of C, Ne, Mg, Si and S with respect to Fe suggest that the interstellar medium in the nucleus is enriched and heated by type-II supernova explosions and winds from massive stars. The cumulative luminosity distribution of the discrete X-ray sources is neither a single nor a broken power law. Separating the sources in the nuclear region (within a distance of 60\\arcsec~from the X-ray centre) from the rest reveals that the two groups have different luminosity distributions. The log~N($>$S) -- log S curve of the sources in the inner region (nucleus and stellar bar) is a single power law, which we interpret as due to continuous, ongoing star formation. Outside the central region, there is a smaller fraction of sources brighter than the Eddington limit for an accreting neutron star. ",
        "introduction": "M83 (NGC~5236) is a grand-design, barred spiral galaxy (Hubble type SAB(s)c) with a starburst nucleus. Distance estimates are still very uncertain. A value of 3.7~Mpc was obtained by de Vaucouleurs et al.\\ (1991). This places the galaxy in the Centaurus A group, whose members have a large spread in morphology and high velocities, indicating that the group is not virialised and tidal interactions and merging are frequent (de Vaucoulers 1979; C{\\^ o}t{\\' e} et al.\\ 1997). A distance of 8.9~Mpc was instead given in Sandage \\& Tamman (1987).  Infra-red (IR) observations (Gallais et al.\\ 1991; Elmegreen, Chromey \\& Warren 1998; Thatte, Tecza \\& Genzel 2000) have shown that the nuclear region of M83 has a complex structure. From measurements of line-of-sight stellar velocities in the inner galactic region, two dynamical centres are inferred. The first centre is identified with a strong point-like optical/IR source. The second centre, located $1\\farcs5$ to the South and $3\\farcs0$ to the West of the IR peak, is not associated with any bright source, but is approximately coincident with the centre of symmetry of the outer isophote ellipses. The stellar velocity dispersion implies that each dynamical centre contains an enclosed mass of $\\approx 1.3 \\times 10^{7}$ M$_{\\odot}$ (Thatte et al.\\  2000).  The $J-K$ images of the nuclear region (Elmegreen et al.\\ 1998) show two circumnuclear dust rings. The inner one has a radius of $2\\farcs8$ and is centred on the IR nuclear peak; the outer one has a radius of $8\\farcs6$ and is centred $2\\farcs5$ South-West of the IR nucleus. The two rings are connected by a mini bar, oriented almost perpendicularly to the main galactic bar. Starburst activity is concentrated in a semi-circular annulus located $\\approx 7$\\arcsec~South-West of the IR nucleus, just inside the South-West half of the outer dust ring (Elmegreen et al.\\ 1998; Harris et al.\\ 2001).  M83 was observed in the X-ray bands by {\\em Einstein} in 1979--1981 (Trinchieri, Fabbiano \\& Palumbo 1985), by {\\em ROSAT} in 1992--1994, and by {\\em ASCA} in 1994 (Okada, Mitsuda \\& Dotani 1997). Thirteen point sources were found in the {\\em ROSAT} PSPC image (Ehle et al.\\ 1998) and 37 in the {\\em ROSAT} HRI image. After removing probable foreground stars and background AGN, there are 21 sources within the D$_{25}$ ellipse believed to belong to the galaxy (Immler et al.\\ 1999). The nuclear region, which encloses approximately 25\\% of the total X-ray luminosity in the 0.1--2.4~keV band, was unresolved in the {\\em ROSAT} images.  M83 was observed by {\\em Chandra} in 2000 April, and the data became available to the public in mid-2001. In this paper we report the results of our analysis of the {\\em Chandra} archival data. We discuss the source population in the galaxy and the properties of discrete sources and unresolved emission in the nuclear region.    \\begin{figure} \\vspace*{0.25cm} \\psfig{figure=rosat_chandra_vlt.ps,width=6.5cm,angle=270} \\caption{ Spatial distribution of the {\\em Chandra} sources (red circles) detected in the S3 chip (green square), and of the {\\em ROSAT} HRI sources (blue crosses with sizes denoting the positional uncertainty). The central region delimited by a dashed box is shown in greater detail in Figure~2. } \\end{figure} ",
        "conclusions": "\\subsection{Luminosity distribution of the sources}  Recent observations have shown that the luminosity distributions of X-ray sources in nearby galaxies can often be approximated by a power-law (eg, disk sources in M101, Pence et al.\\ 2001, and M81, Tennant et al.\\ 2001) or by a broken power-law profile (eg, bulge sources in M31, Shirey et al.\\ 2001, and in M81, Tennant et al.\\ 2001). The log~N($>$S) -- log~S curves for the sources located in the bulges of spiral galaxies tend to have steeper slopes at the high-luminosity end. For example, the broken power-law distribution for bulge sources in M31 has slopes of $-0.5$ and $-1.8$ at the low- and high-luminosity end, respectively (Shirey et al.\\ 2001). The log~N($>$S) -- log~S curves for the sources in galactic disks are generally single, flatter power laws (eg, with a slope of $-0.5$ for the disk sources in M81, Tennant et al.\\ 2001, and $-0.8$ for the disk sources in M101, Pence et al.\\ 2001). Elliptical galaxies also have broken power-law log~N($>$S) -- log~S curves, with generally steep slopes at their bright ends (eg, with a slope of $-1.8$ for NGC~4697, Sarazin, Irwin \\& Bregman 2000), similar to those inferred for the bulge sources in spiral galaxies. Starburst galaxies, instead, tend to have flat power-law log~N($>$S) -- log~S curves (eg, with a slope of $-0.45$ for NGC~4038, Fabbiano, Zezas \\& Murray 2001) similar to the distributions for the disk sources in spiral galaxies.  A flatter power-law luminosity distribution implies a larger proportion of bright sources in a population. If the brightest X-ray sources in a galaxy are young, short-lived high-mass X-ray binaries, born in a recent starburst episode, then the slope of the bright end of the log~N($>$S) -- log~S curve indicates the star-formation activity of the galaxy (see Prestwich 2001).  Various mechanisms can produce a broken power-law profile in the log~N($>$S) -- log~S curve. The break may be caused by a pile-up of systems at a particular luminosity; in particular, it may be due to a population of neutron-star X-ray binaries (Sarazin, Irwin \\& Bregman 2000) with a mass-transfer rate at or just above the Eddington limit, i.e., with bolometric luminosities $\\sim 2 \\times 10^{38}$~erg~s$^{-1}$. Aging of a population of X-ray binaries born during a starburst episode can also produce a luminosity break in the log~N($>$S) -- log~S curve (Wu 2001; Wu et al.\\ 2001). In this model, the initial distribution has a power-law profile; as the bright, short-lived systems die, a break is created, moving gradually to lower luminosities with time. Pile-up of neutron-star X-ray binaries is a likely cause of the luminosity break in the log~N($>$S) -- log~S curves of elliptical galaxies, where active star formation is absent. Population aging is a more likely mechanism for spiral galaxies, in particular those that have had tidal interactions with their satellites in the recent past (Wu 2001; Wu et al.\\ 2001).    We have shown in Section 3 that the sources in the nuclear region of M83 and those in the disk appear to have different luminosity distributions. One of the obvious features in the log~N($>$S) -- log~S curve of the disk sources is a kink, located roughly at the Eddington luminosity of accreting neutron stars (if a distance of 3.7~Mpc is assumed). The slopes are approximately $-0.6$ and $-1.3$ at the low- and high-luminosity end, respectively. When we consider only the sources inside the 60\\arcsec~central region (which includes the nucleus and the stellar bar, but not the spiral arms), we obtain a simple power-law distribution with a slope of $-0.8$ (Figure 3). This implies that the population of sources in the inner regions (nucleus and bar) has a larger relative fraction of bright sources than the disk population.  The situation is different for example in the spiral galaxy M81, where most bright sources are found in the galactic disk instead of the nuclear region (Tennant et al.\\ 2001). If the flatness of the slope in the log~N($>$S) -- log~S curve is an indicator of recent star formation, the difference in the spatial distribution of the brightest sources in M83 and M81 is simply a consequence of the fact that M83 has a starburst nucleus while star formation in galaxies such as M81 is presently more efficient in the galactic disk. With this interpretation, the current star formation rate in the disk of M83 would be intermediate between the rate in the disk of M81 (slope of $-0.5$ at the high-luminosity end) and in the bulge of M31 (slope of $-1.8$).  \\subsection{Nuclear region}  The distribution of bright, young star clusters suggests that the most vigorous star formation is concentrated in a semi-circular annulus $\\simeq 7\\arcsec$~($\\approx 130$ pc) South-West of the IR photometric nucleus. From the colour distribution of the star clusters, it is inferred that star formation first started at the southern end of the ringlet about $10$--$30$~Myr~ago (Thatte et al.\\ 2000; Harris et al.\\ 2001), and has since propagated towards the currently more active northern end, where the age of most young stellar clusters is $\\simlt 5$~Myr. The youngest clusters are found along the outer edge of the ringlet, indicating that the star-formation front is presently propagating outwards (Harris et al.\\ 2001).  The starburst nucleus of M83 was unresolved in all X-ray observations before {\\em Chandra}; therefore, no comparisons between observations in the X-rays and in other wavelengths had been possible. Overplotting the {\\em Chandra} brightness contours in the 0.3--8.0~keV band on the {\\em HST}/WFPC2 multi-colour images (Figure~11) helps to shed light on the relative spatial distribution of the X-ray and optical emission. For instance, although there is a general correlation, the optical emission is more strongly concentrated around the IR photometric nucleus and along the star-forming ring, while the X-ray emission is more uniformly distributed. There is also extended X-ray emission South-West of the star-forming arc, along the direction of the main galactic bar, and towards the North-West, across the dust lane, in a direction perpendicular to the bar. Discrete X-ray sources and unresolved emission are also found to the East of the IR photometric nucleus, not associated with any bright optical regions with currently active star formation. This may be due to the fact that the optical emission traces the bright young stellar clusters and OB associations, while the X-ray emission is in general associated with remnants of stellar evolution such as accreting compact stars and supernovae.  The HST image shows that the IR photometric nucleus is strongly extincted in the UV. The X-ray spectrum of the corresponding {\\it Chandra} source (No.~43 in Table A.1) does not allow a precise determination of the absorption column density, but it constrains it to be $< 2.4 \\times 10^{21}$~cm$^{-2}$ (Table 3). From the relation between the absorption column density $n_{\\rm H}$ and the visual extinction $A_V$ (Predehl \\& Schmitt 1995), this implies that the source has a visual extinction $A_V < 1.3$ mag. An extinction $A_V = 0.9$ mag has been deduced from IR observations (Thatte et al.\\ 2000).  Unresolved soft emission extends for $\\simgt 15$\\arcsec~ ($\\simgt 270$~pc) to the South-West, West and North-West, outside the star-forming arc; the unresolved hard X-ray emission appears instead to be confined in the region between the IR photometric nucleus and the star-forming arc (Figure 4, bottom right panel), and may perhaps extend to the North-West for $\\simlt 10$\\arcsec~ ($\\approx 180$ pc). The X-ray spectrum of the unresolved component shows strong emission lines, typical of emission from optically-thin thermal plasma at $kT \\approx 0.6$ keV.  Above-solar abundances of Ne, Mg, Si and S are required to fit the spectrum, while Fe appears to be underabundant. This suggests that the interstellar medium in the starburst nuclear region has been enriched by the ejecta of type-II supernova explosions. A high C abundance and a high C/O abundance ratio can be the effect of radiatively-driven winds from metal-rich massive stars ($M \\simgt 40$\\,M$_{\\odot}$) in their Wolf-Rayet stage (Gustafsson et al 1998; Portinari, Chiosi \\& Bressan 1998). Both effects are likely to be present in the nuclear region.  Four discrete sources are resolved along the star-forming arc. The brightest of them (No.~44 in Table A.1) is located near the southern end of the arc, where star formation started $\\approx 10$--$30$~Myr~ago (Harris et al.\\ 2001). The other three sources (No.~35, 37 and 40) are towards the northern end of the arc, where star formation started more recently ($\\simlt 5$ Myr ago); source No.~35 is approximately coincident with the location of SN~1968L, a type-II supernova (Wood \\& Andrews 1974).     \\begin{figure*} \\vbox{ \\begin{tabular}{lr} \\psfig{figure=pc300_xv3.ps,width=8.5cm} & \\psfig{figure=pc547_xv3.ps,width=8.5cm} \\\\ \\psfig{figure=pc814_xv3.ps,width=8.5cm} & \\psfig{figure=hst_color3_xv3.ps,width=8.35cm} \\\\ \\end{tabular} } \\caption{The contours of the 0.3--8.0 keV emission in the nuclear region (smoothed with a $1\\farcs5 \\times 1\\farcs5$ boxcar kernel) are overplotted on the {\\em HST}/WFPC2 images in the F300W ($\\approx U$; top left panel), F547M ($\\approx V$; top right) and F814W ($\\approx I$; bottom left) broad-band filters. A three-colour {\\em HST}/WFPC2 image is also shown for comparison (bottom right, from Harris et al.\\ 2001). The HST images, taken on 2000 April--May, were obtained from the STScI public archive. North is up, East is left. } \\label{fig:image} \\end{figure*}  \\subsection{Black-hole candidates}  Active galactic nuclei (AGN) are often found to have power-law X-ray spectra. If a thermal blackbody component is present, it peaks at energies $\\simlt 0.1$ keV and does not dominate in the 0.3--8.0~keV band. Black-hole X-ray binaries in hard spectral states also have single power-law spectra in the 0.3--8.0~keV band. When they are in the soft spectral state, instead, their X-ray spectra consist of a strong thermal blackbody component, with a temperature of about 1~keV, and a power-law tail, usually steeper than the power law in the hard state.   The IR photometric nucleus has a power-law X-ray spectrum consistent with either those of accreting supermassive black holes in AGN, or those of stellar-mass black-hole candidates in X-ray binaries, in the hard state. The mass associated to the source is estimated, from stellar kinematics, to be $\\approx 10^7$~M$_{\\odot}$ (Thatte et al.\\ 2000). If the source contains a supermassive black hole, analogous to those found in the nuclei of common Seyferts and quasars, and if the X-ray emission is powered by accretion, its luminosity ($L_{\\rm x} \\simeq 3.2 \\times 10^{38}$~erg~s$^{-1}$ in the 0.3--8.0~keV band) is well below the Eddington limit; therefore, M83 is not in an AGN phase in the present epoch.  In the most accepted scenario, short-lived massive stars are the progenitors of stellar-mass black-hole candidates. Thus, some of the discrete X-ray sources in the nuclear region (excluding possible supernova remnants and the IR nucleus itself) may be high-mass black-hole X-ray binaries. Source No.~44 is a likely black-hole candidate, end product of the most recent starburst episode in the region (age $< 30$ Myr): its unabsorbed luminosity $\\simgt 2.4 \\times 10^{38}$~erg~s$^{-1}$ in the 0.3--8.0~keV band places it above the luminosity of a neutron star accreting material at the Eddington rate.  Although the two brightest nuclear sources, No.~43 and 44, have comparable luminosity, and they are probably both powered by accretion onto a black hole, their spectral properties are different. The hard power-law photon index ($\\Gamma < 1.5$) in the spectrum of the photometric nucleus and the kinematic properties of the stars around it suggest that it is a supermassive black hole. The other source has a much softer power-law photon index ($\\Gamma > 2.5$) consistent with the spectrum of a stellar-mass black-hole candidate in the soft state. The present data are not sufficient to determine whether or not a disk-blackbody component is also present.  We can compare these two nuclear black-hole candidates with the brightest X-ray source in the M83 field, the {\\em ROSAT} HRI source H30, which is located $\\approx 4\\farcm5$ ($\\sim 5$ kpc) away from the nucleus. Immler et al.\\ (1999) suggested that H30 is also a strong black-hole candidate, because its luminosity is well above the Eddington limit for a 1.5-M$_\\odot$ compact star. (The source is still super-Eddington after accounting for the smaller distance to M83 assumed in this work.) Its spectrum can be fitted equally well with various models (Table 5). The most likely physical interpretation is that it is dominated by a steep ($\\Gamma \\approx 2.5$) power-law component, produced by Comptonisation of a soft blackbody component. The current data do not allow a simultaneous determination of the absorbing column density and of the temperature of the blackbody component. If the blackbody temperature $T_{\\rm bb} \\ll 0.2$ keV, the total column density $n_{\\rm H} \\approx 2.5 \\times 10^{21}$~cm$^{-2}$. In the other extreme case, for $n_{\\rm H} = 4 \\times 10^{20}$~cm$^{-2}$ (no intrinsic absorption), $T_{\\rm bb} \\approx 0.25$ keV. An accurate determination of the optical extinction of a possible optical counterpart would constrain the temperature of the blackbody component in the X-ray spectrum, and hence the mass of the accreting object. This is left to further work.  The spectrum of H30 can also be well fitted with an optically-thin thermal plasma model (bremsstrahlung emission from completely ionised plasma at $kT \\approx 2.5$ keV), consistent with a luminous supernova remnant. However, this latter possibility is much more unlikely because the source has shown variability during the {\\em ROSAT} HRI observations and over the timespan of the {\\em Chandra} observations. Moreover, no X-ray emission lines are detected in the spectrum, but the best-fit temperature is too low for the plasma to be completely collisionally-ionised. This would imply a metal abundance $\\simlt 0.1$ times the solar value (we have estimated it by fitting an absorbed Raymond-Smith model in {\\small {XSPEC}}). Alternatively, the gas could be photo-ionised. Finally, a simple disk-blackbody spectrum with colour temperature of the innner accretion disk $T_{\\rm in} \\approx 0.9$ keV, although physically unlikely, cannot be ruled out with the current data."
    },
    "0201/astro-ph0201090_arXiv.txt": {
        "abstract": "We calculate chemical evolution models for 4 dwarf spheroidal satellites of the Milky Way (Carina, Ursa Minor, Leo I and Leo II) for which reliable non-parametric star formation histories have been derived. In this way, the independently obtained star formation histories are used to constrain the evolution of the systems we are treating. This allows us to obtain robust inferences on the history of such crucial parameters of galactic evolution as gas infall, gas outflows and global metallicities for these systems. We can then trace the metallicity and abundance ratios of the stars formed, the gas present at any time within the systems and the details of gas ejection, of relevance to enrichment of the galaxies environment. We find that galaxies showing one single burst of star formation (Ursa Minor and Leo II) require a dark halo slightly larger that the current estimates for their tidal radii, or the presence of a metal rich selective wind, which might carry away much of the energy output of their supernovae before this might have interacted and heated the gas content, for the gas to be retained until the observed stellar population have formed. Systems showing extended star formation histories (Carina and Leo I) however, are consistent with the idea that their tidally limited dark haloes provide the necessary gravitational potential wells to retain their gas. The complex time structure of the star formation in these systems, remains difficult to understand. Observations of detailed abundance ratios for Ursa Minor, strongly suggest that the star formation history of this galaxy might in fact resemble the complex picture presented by Carina or Leo I, but localized at a very early epoch. ",
        "introduction": "Our understanding of the dwarf spheroidal companions of the Milky Way has advanced significantly over the last few years. Studies of the internal dynamics of stars measuring velocity dispersions, have revealed the presence of varying but always significant amounts of dark matter (e.g. Mateo et al. 1993, Tamura et al. 2001), strengthening the classification of these systems as galaxies. The structure of the dark haloes associated with these galaxies appears to be well represented by a constant density region, over the extent across which measurements exist. High quality imaging of the stellar populations has permitted the construction of HR diagrams for many of these galaxies, which in turn has stimulated the development of new statistical analysis techniques aimed at recovering the star formation histories of the imaged systems e.g. Aparicio \\& Gallart (1995), Tolstoy \\& Saha (1996), Mighell (1997), Hernandez et al. (1999). These last have yielded valuable constraints on the temporal structure of the star formation histories in the local dSph population. Results have shown there to be a wide range of different star formation histories here. Some systems being consistent with a single burst of activity in the remote past, but many showing clear signs of extended or repeated periods of star formation activity.   The dominance of dwarf galaxies in cluster luminosity functions indicates that the collective contribution of these systems to the chemical evolution of clusters is quite probably significant. Detailed studies of the role of small galactic systems in the context of the enrichment of the inter-galactic medium in clusters have confirmed the above, generally treating dwarf and dSph galaxies as small scale analogues of elliptical galaxies, typically characterized by a single burst of star formation in the remote past, and nothing since e.g. Gibson \\& Matteucci (1997). Recently, the prevalence of higher than primordial metal abundances in the high redshift $Ly\\alpha$ systems was taken by Madau, Ferrara \\& Rees (2001) as an unequivocal indication of substantial metal pollution by galactic systems of masses in the dSph range at high redshift. The details of the star formation physics and efficiency in these small systems remains as a free parameter which influences the details of their calculations significantly, again approximated by a single burst of activity. However, given the complex star formation histories local dSph's present, it is worthwhile to analyze more closely what a careful tracing of the chemistry in these systems, fixed so as to reproduce their observed star formation histories within their observed dark haloes, predicts in terms of metallicities, abundance ratios and total amounts of ejected gas. The smallest galactic structures also play a determinant part in models of reionization of the universe, where again, these systems are usually treated only in a very generic way (e.g. Ciardi, Ferrara \\& Abel 2000). The local dSphs treated in this work offer a unique opportunity to study the details of galactic evolution at the smallest scales, and hence calibrate the details of their contribution to a variety of astrophysical problems.  The observation of complex stellar populations in the dSph satellites of the Milky Way, in the total absence of gas (e.g. Bowen et al. 1997), implies that the gas which was at some point present in these systems to fuel star formation processes, has now been totally lost. Given their dynamical masses, theoretical models of gas heating due to supernovae (henceforth SN) and stellar winds in these systems predict the efficient formation of galactic winds after only a few hundred Myrs of star forming activity, with the resulting loss of all gas (e.g. Dekel \\& Silk 1986, Ferrara \\& Tolstoy 2000). As already pointed out by Gallagher \\& Wyse (1994), the low escape velocities of dSph's implies that the evolution of gaseous and stellar metallicities will be determined mostly by the inflows and outflows of gas, rather than by the details of the star formation history, as is the case in larger systems, where the use of closed box models is generally accepted as a good first approximation. Note however, the recent detection of some HI in Sculptor by Carignan et al. (1998), a galaxy with some indication of a young stellar component superimposed on a dominant old population. The distance to Sculptor does not allow the detailed study of its stellar population, so that the inferences of its SFH are only tentative, and doubts remain concerning the actual physical association of the observed gas to the galaxy. This again stresses the poor knowledge of the interstellar medium-star formation rate (henceforth ISM and SFR, respectively) connection in these systems, a consequence of the absence of a detailed physics for star formation.  The presence of stellar populations showing age spreads of several Gyrs, remains the greatest puzzle surrounding these small galaxies. Re-accretion of gas has been proposed as the only plausible explanation for the extended star formation histories observed, although the details and causes of any repeated periods of accretion also remain a mystery. In this work we aim at restricting the parameters describing repeated periods of gas accretion in some of the local dSph galaxies, as well as obtaining indications as to the evolution of the metallicity in them.  We construct a chemical model for the evolution of the stellar and gaseous constituents of a small galactic system, paying attention to the enrichment processes, and the heating and dynamics of the gas. The dynamics being largely determined by the dark halo of the systems, which we fix from independent inferences. The determinant input of the model is now the accretion history of fresh gas, with the metallicity of the stars and gas, together with the appearance of a galactic wind being the output of the model. In absence of a detailed formation and evolution scenario, this gas accretion history must be treated through free parameters. Fortunately for 4 of the local dSph's, direct and objective determinations of their star formation histories exist. Hernandez, Gilmore \\& Valls-Gabaud (2000), henceforth HGC, use a rigorous maximum likelihood statistical method to derive the SFR(t) for 4 dSph galaxies, Carina, Ursa Minor, Leo I and Leo II, in a totally non parametric approach.  In a complementary approach, Carraro et al. (2001) model dSph systems from first principles, and show that for systems having a fixed mass, in the dSph range, variations in the total dynamical density can result in a range of different star formation histories, spanning that presented by the satellites of the Milky Way. Again, the star formation process was introduced through dimensional arguments and free parameters, in the absence of detailed knowledge of star formation physics, and the galaxies are treated as isolated systems. This last assumption probably does not hold for local dSph's where the presence of the MW exerts a possibly dominant effect on the evolution. The influence of the tidal field of the Galaxy is well established in several of these systems, Martinez-Delgado et al. (2001) find evidence of tidal tails in Carina, and Majewski et al. (2000) determine the presence of similar streams of drawn out material in Ursa Minor. Further, Hirashita, Kamaya \\& Mineshige (1997) point to the dynamical frailty of the gas in dSph systems in relation to the gas loss associated to the hot outflows produced by the star formation processes in the disk of the Milky Way. Scanapieco, Ferrara \\& Broadhurst (2000) confirm the above through detailed dynamical modeling of the photo evaporation and ram striping of gas from dSph's due to galactic winds and fountains. All this only stresses the fact that these deceivingly simple galactic systems are subject to complex processes which make it difficult to construct physical models to describe their evolution.   In our present study we take the SFR's derived by HGC (henceforth $SFR_{HGV}$) as external constraints on our chemical evolution models, and hence obtain interesting restrictions on the time structure and magnitude of the gas accretion history of Carina, Ursa Minor, Leo I and Leo II. In a sense, for these galaxies we know part of the answer in advance, and can hence calculate the energy input produced by the inferred star formation history, and restrict the possible gas accretion and outflow scenarios, thus obtaining valuable information on the nature of the ISM-SFR connection. The time evolution of the metallicities is then a prediction of the model, which we can compare with observed values.  The models we obtain show a variety of possibilities for the physical evolution of these systems, depending on which parameters one varies to ensure gas is retained until the luminous galaxy is formed. Galaxies showing repeated periods of star formation, such as Carina and Leo I in our sample, can only be explained with the inclusion of a re-accretion of fresh gas. We obtain predictions on the total masses, metallicities and abundance ratios of the ejected material, as a result of having carefully traced the physics of the gas content, the different SN yields, and the final results of the outflows. A simple physical criterion is also proposed as relevant to discriminating dSph galaxies subject to extended and repeated star formation, from those susceptible only to a single burst of activity.   The plan of our paper is as follows: Section 2 presents the details of the enrichment and gas dynamics model, with the results once the inferred SFR(t)'s have been introduced as constraints, presented in Section 3. Finally, a discussion of our results is given in Section 4 and Section 5 states our conclusions. ",
        "conclusions": "To conclude:  1) We have performed a detailed exploration of the parameter space available to the four local dSph galaxies studied here, in terms of the gas accretion regimes and metallicities of the formed stars and ejected materials, by taking as external restrictions the star formation histories of these systems, as inferred from direct statistical studies of their resolved populations, together with the observed general properties of their dark haloes. This shows that knowledge of the time structure and normalization of the star formation rate of external galaxies can be combined with physical and chemical modeling of these systems to derive interesting information on their evolution and dark matter haloes, the details of which furnish the boundary conditions within which the galaxies evolve.  2) From the observed abundance ratios of Ursa Minor, in combination with the physics of gas flows and our chemical models, we find strong suggestions that this galaxy, often thought of as the prototypical old, single burst and low metallicity dSph, might in fact have experienced a complex star formation history. This has remained hidden from direct studies of its resolved stellar population by its large age, i.e. everything happened more than 10 Gyrs ago, but it was not simple.  3) We found evidence for a slight underestimate of $R_{tidal}$ as the total extent of their dark halos or of a significant metal rich ejecta, with reality probably falling somewhere in between.  4) Comparison of the predicted abundance ratios with the available data shows a broad consistency of our chemical and physical modeling with the relevant observations.  5) A simple physical criterion to estimate when a dSph system might sustain extended star formation, as opposed to being subject to a single burst of activity is presented, which neatly separates into those two classes the galaxies we studied."
    },
    "0201/astro-ph0201084_arXiv.txt": {
        "abstract": "{The time variation in the water-vapour bands in oxygen-rich Mira variables has been investigated using multi-epoch ISO/SWS spectra of four Mira variables in the 2.5--4.0\\,$\\mu$m region. All four stars show H$_2$O bands in absorption around minimum in the visual light curve. At maximum, H$_2$O emission features appear in the $\\sim$3.5--4.0\\,$\\mu$m region, while the features at shorter wavelengths remain in absorption. These H$_2$O bands in the 2.5--4.0\\,$\\mu$m region originate from the extended atmosphere.  The analysis has been carried out with a disk shape, slab geometry model. The observed H$_2$O bands are reproduced by two layers; a `hot' layer with an excitation temperature of 2000\\,K and a `cool' layer with an excitation temperature of 1000--1400\\,K. The column densities of the `hot' layer are $6\\times10^{20}$--$3\\times10^{22}$\\,cm$^{-2}$, and exceed $3\\times10^{21}$\\,cm$^{-2}$ when the features are observed in emission. The radii of the `hot' layer (\\Rhot) are $\\sim$1\\,$R_*$ at visual minimum and 2\\,$R_*$ at maximum, where $R_*$ is a radius of background source of the model, in practical, the radius of a 3000\\,K black body. The `cool' layer has the column density (\\Ncool) of $7\\times10^{20}$--$5\\times10^{22}$\\,cm$^{-2}$, and is located at 2.5--4.0\\,$R_*$. \\Ncool\\ depends on the object rather than the variability phase.  The time variation of $\\Rhot/R_*$ from 1 to 2 is attributed to the actual variation in the radius of the H$_2$O layer, since the variation in \\Rhot\\ far exceeds the variation in the `continuum' stellar radius. A high H$_2$O density shell occurs near the surface of the star around minimum, and moves out with the stellar pulsation. This shell gradually fades away after maximum, and a new high H$_2$O density shell is formed in the inner region again at the next minimum. Due to large optical depth of H$_2$O, the near-infrared variability is dominated by the H$_2$O layer, and the L'-band flux correlates with the area of the H$_2$O shell. The infrared molecular bands trace the structure of the extended atmosphere and impose appreciable effects on near-infrared light curve of Mira variables. ",
        "introduction": "Asymptotic Giant Branch (AGB) stars are in the late stage of the stellar evolution for low and intermediate main-sequence mass stars. Generally, AGB stars are pulsating variables as represented by Mira variables. According to hydrodynamic model atmospheres, the pulsations lift up matter from the stellar surface and extend the atmosphere (e.g. Bowen~\\cite{Bowen88}). Pulsations create shocks, causing a step-like structure in the density distribution as a function of radius (e.g. Fleischer et al.~\\cite{Fleischer92}; H\\\"ofner et al.~\\cite{Hoefner98}). The cooling behind the shock is efficient (Woitke et al.~\\cite{Woitke96}) and the temperature decreases immediately in the post-shock regions (Fleischer et al.~\\cite{Fleischer92}; H\\\"ofner et al.~\\cite{Hoefner98}).  The extended atmosphere is filled with various kinds of molecules. The structure of the extended atmosphere can be studied using infrared molecular bands. Hinkle (\\cite{Hinkle78}) and Hinkle \\& Barnes (\\cite{Hinkle79}) analyzed high resolution spectra of the oxygen-rich Mira, \\object{R~Leo}, and found two velocity components in the near-infrared molecular lines of CO, OH, and H$_2$O. They concluded that one component is located near the boundary region of the photosphere and the second component is superposed on the first layer above the photosphere.  However, the molecular bands suffer interference from molecules in the terrestrial atmosphere. Recent space-borne observations enable more comprehensive studies of the molecules in the extended atmosphere. Using the Short-Wavelength Spectrometer (SWS; de~Graauw et al.~\\cite{deGraauw96}) on board the Infrared Space Observatory (ISO; Kessler et al.~\\cite{Kessler96}), Tsuji et al. (\\cite{Tsuji97}) found CO, H$_2$O, CO$_2$ and SiO molecules located above the photosphere. Markwick \\& Millar (\\cite{Markwick00}) indicated that the 2.8\\,$\\mu$m spectra of a Mira variable consists of two H$_2$O components with 950\\,K and 250\\,K. Yamamura et al. (\\cite{Yamamura99a}) identified SO$_2$ features in the 7\\,$\\mu$m region in three oxygen-rich Mira variables with an excitation temperature estimated to be 600\\,K. The band was found to be variable, changing from emission to absorption. The time scale of the SO$_2$ variation was longer than the period of the visual variable phase. Justtanont et al. (\\cite{Justtanont98}) and Ryde et al. (\\cite{Ryde99a}) detected CO$_2$ bands in the 12--17\\,$\\mu$m region. Cami et al. (\\cite{Cami00}) found these CO$_2$ molecules fill the region between 4--400\\,stellar radii. Not only molecules but also fine-structure atomic lines were found in the upper atmosphere (Aoki et al.~\\cite{Aoki98a}).  H$_2$O is one of the most abundant molecules in the oxygen-rich atmosphere and is a large opacity source in the near-infrared region. Tsuji (\\cite{Tsuji78b}), using low resolution spectra obtained by the Kuiper Airborne Observatory, suggested that the 5--8\\,$\\mu$m region is filled with the H$_2$O emission arising from the atmosphere above the photosphere (or the hydrostatic atmosphere). However, the spectral resolution of those data was too low for further study. Yamamura et al. (\\cite{Yamamura99b}) found emission features from water-vapour bands around 3.5--4.0\\,$\\mu$m in $o$~Cet, which was observed at maximum in the visual light curve. In contrast, \\object{Z~Cas}, which was observed near minimum, showed absorption features at the same wavelengths. They analyzed water-vapour spectra from 2.5 to 4.0\\,$\\mu$m with a simple `slab' model. The model consists of two molecular layers (`hot' layer and `cool' layer) with independent excitation temperatures, column densities, and radii. The `hot' layer with an excitation temperature of 2000\\,K extended to $\\sim$2\\,$R_*$ in $o$~Cet and stayed at $\\sim$1\\,$R_*$ in Z~Cas, where $R_*$ is the radius of the background light source representing the star. They surmised that water layers are generally more extended at maximum.  In this paper we examine the time variation in the water vapour bands in the 2.5--4.0\\,$\\mu$m region. Yamamura et al. (\\cite{Yamamura99b}) analyzed two different stars at different phases. To investigate whether the difference in the radii of the water layers is related to the phase difference or not, we analyzed the spectra of four Mira variables observed several times with the ISO/SWS. We found periodical variation in H$_2$O bands in ISO/SWS spectra. Some features turn from absorption to emission during minimum and maximum. This variation is explained by the variation in the radius of the H$_2$O layer. These H$_2$O molecules are located in the extended atmosphere. We discuss the variation in the structure of the extended atmosphere caused by the pulsations. ",
        "conclusions": "We reported the time variation in H$_2$O bands in four oxygen-rich Mira variables. The near infrared water-vapour bands at 2.5--3.95\\,$\\mu$m follow the periodical variation. Emission features are seen at $\\sim$3.5--3.95\\,$\\mu$m around maximum while absorption features are detected around minimum. The spectra are well fitted with `slab' models, which consist of two H$_2$O layers (hot layer and cool layer). The radius of the hot layer varies from $\\sim$1\\,$R_*$ to $\\sim$2\\,$R_*$ during visual minimum and maximum.  The periodical variation in the features arising from the extended atmosphere suggests that the structure of the outer atmosphere is varying with the pulsation. The pulsation produces a shock in the atmosphere, and the hot H$_2$O layer traces the high H$_2$O density caused by the shock. The high H$_2$O density shell expands from inside of the extended atmosphere to outwards from minimum to maximum. In our analysis, $R_*$ is measured as the relative number of the stellar radius. Considering the variation in the radius of the star, the high density shell is by a factor of 1.5--2.0 more extended at maximum than at minimum.  Due to large optical depth of H$_2$O bands, the spectra in the 2.5--4.0\\,$\\mu$m region are dominated by the H$_2$O in the extended atmosphere. The flux variation in the L'-band is primarily determined by the radial motion of optically thick H$_2$O layer."
    },
    "0201/astro-ph0201421_arXiv.txt": {
        "abstract": "Observations of the $X$-ray band wavelength reveal an evident ellipticity of many galaxy clusters atmospheres. The modeling of the intracluster gas with an ellipsoidal $\\beta$-model leads to different estimates for the total gravitational mass and the gas mass fraction of the cluster than those one finds for a spherical $\\beta$-model. An analysis of a recent {\\it Chandra} image of the galaxy cluster RBS797 indicates a strong ellipticity and thus a pronounced aspherical geometry. A preliminary investigation which takes into account an ellipsoidal shape for this cluster gives different mass estimates than by assuming spherical symmetry. We have also investigated the influence of aspherical geometries of galaxy clusters, and of polytropic profiles of the temperature on the estimate of the Hubble constant through the Sunyaev-Zel'dovich (SZ) effect. We find that the non-inclusion of such effects can induce errors up to 40 \\% on the Hubble constant value. ",
        "introduction": "Clusters of galaxies can be used to study how structures form on large scales. The formation and evolution of clusters depend very sensitively on cosmological parameters like the mean matter density $\\Omega_m$ in the Universe. Thus it is of great importance to determine the dynamical state of clusters at different redshifts, see e.g. \\cite{sabi2001}. \\\\ Geometry can give important insight in the dynamics of galaxy clusters. For example, the fitting of the cluster $X$-ray surface brightness with $\\beta$-models usually provides different best fit parameters depending on the shape one assumes for the intracluster gas, but generally  the {\\it classical} calculations of the mass of galaxy clusters suppose a spherical distribution of the density. \\\\ The SZ effect, \\cite{zel1972}, offers the possibility to put important constraints on the cosmological models. Combining the temperature change in the cosmic microwave background due to the SZ effect and the $X$-ray emission observations, the angular distance to galaxy clusters and thus the Hubble constant $H_o$ can be derived. Nevertheless, geometrical shape of galaxy clusters can also introduce some {\\it errors} on the analysis of the SZ effect. \\\\ In Section 2 we will analyze some consequences of the geometry of the clusters of galaxies and particularly on the mass, by taking as an example a recent observation of the galaxy cluster RBS797. Then we will describe, in Section 3, the influence of the cluster shape on the estimate of the SZ effect. In section 4 we will give an outlook on the role played by the geometry of galaxy clusters. ",
        "conclusions": ""
    },
    "0201/astro-ph0201492_arXiv.txt": {
        "abstract": "Recently, \\citet{l02} have measured the distribution of star formation rate intensity in galaxies at various redshifts. This data set has a number of advantages relative to galaxy luminosity functions; the effect of surface-brightness dimming on the selection function is simpler to understand, and this data set also probes the size distribution of galactic disks. We predict this function using semi-analytic models of hierarchical galaxy formation in a $\\Lambda$CDM cosmology. We show that the basic trends found in the data follow naturally from the redshift evolution of dark matter halos. The data are consistent with a constant efficiency of turning gas into stars in galaxies, with a best-fit value of $2\\%$, where dust obscuration is neglected; equivalently, the data are consistent with a cosmic star formation rate which is constant to within a factor of two at all redshifts above two. However, the practical ability to use this kind of distribution to measure the total cosmic star formation rate is limited by the predicted shape of an approximate power law with a smoothly varying power, without a sharp break. ",
        "introduction": "\\label{intro}  One of the major goals of the study of galaxy formation is to achieve an observational determination and a theoretical understanding of the cosmic star formation history. Previous measurements have generally found a factor of $\\sim 4-10$ increase in the cosmic star formation rate (henceforth SFR) from redshift $z=0$ out to $z \\sim 1$--2, with the cosmic SFR remaining roughly constant at higher redshifts out to $z \\sim 4$, once approximate corrections are included for incompleteness or for the effect of dust extinction.  In general, estimates of the SFR apply locally-calibrated correlations between emission in particular lines or wavebands and the total SFR. The observational picture is based on a large number of observations in different wavebands. These include various ultraviolet/optical/near-infrared observations \\citep[e.g.,] []{m96,l96,mpd98,t98,p98,csb99,g99,f99,s99,hcs00}. At the shortest wavelengths, the extinction correction may be large (a factor of $\\sim 5$ at the highest redshifts) and is still highly uncertain. At longer wavelengths, the star formation history has been reconstructed from submillimeter observations \\citep{b99,h98} and radio observations \\citep{c98}; in this range of the spectrum, large uncertainties remain because of the insufficiency of current observational constraints on the spectral shape of the galaxies' dust emission.  Hierarchical models have been used in many papers to match observations on star formation at $z \\la 4$ \\citep[e.g.,][] {b98,kc98,spf01}. This comparison should become increasingly profitable for our physical understanding of galaxy formation as observations probe the galaxy population more completely at low redshift and also push towards high redshift.  Recently, \\citet{l02} introduced a new data set, which takes advantage of the deepest images taken by the Hubble Space Telescope (HST), namely the Hubble Deep Field (HDF), and the Hubble Deep Field South (HDF-S) WFPC2 and NICMOS fields. These fields contain a fairly large number of galaxies, including a good fraction at high redshift. Combining a large variety of space- and ground-based optical and infrared images of these fields, \\citet{l02} found photometric redshifts for $\\sim 3000$ galaxies. They measured redshifts using a sequence of six spectral templates which account for different galaxy types. Their redshift technique has been checked with spectroscopic measurements at $z<6$, yielding a relative root-mean-square dispersion of 0.065 in $1+z$; however, the check applies mainly to $z=0$--1.2 and $z=2.2$-3.5, with only a small number of checks outside these redshift ranges \\citep{y00}. Another limitation of the current data is that the fields are relatively small, and are not randomly selected.  The HDF was chosen to be in a relatively empty field, while the HDF-S was chosen to be near a $z=2.2$ quasar. Thus, cosmic variance may be significant, although it should be greatly suppressed by the large redshift bins that are used in the analysis.  The data analysis of \\citet{l02} presents several novelties compared to optical and infrared determinations of the galaxy luminosity function. Since most galaxies are well-resolved in the HST images, \\citet{l02} divide each galaxy into individual pixels, and measure the SFR intensity $x$ in each pixel, in units of $M_{\\odot}$ yr$^{-1}$ kpc$^{-2}$, where all quantities are proper. They then add up the proper areas of all pixels, within a given redshift range, with SFR intensity in the interval $x$ to $x+dx$. Dividing by the comoving volume in the redshift bin, and by $dx$, they obtain $h(x)$, the SFR intensity distribution function, expressed in units of proper kpc$^2$ per comoving Mpc$^3$ per unit of $x$, i.e., kpc$^2$ Mpc$^{-3}/ (M_{\\odot}$ yr$^{-1}$ kpc$^{-2})$. Once $h(x)$ has been obtained, the total cosmic SFR per comoving volume is simply $\\int x\\, h(x)\\, dx$.  Measuring the SFR in pixels offers a number of advantages compared to measuring the total SFR in a galaxy. First, as noted by \\citet{l02}, the total luminosity of a galaxy is in practice not very well defined, because the luminosity is integrated out to a radius where the surface brightness drops below the noise, and this radius depends strongly on redshift due to cosmological surface brightness dimming. In addition, the effect of surface brightness limits on the selection function is non-trivial in the case of galaxy luminosity functions; specifically, whether or not a galaxy is detected does not depend only on its luminosity, but also on its size and its orientation relative to the line of sight. On the other hand, since each pixel has a definite, known angular size, the selection function is simple, i.e., there is a minimum $x$ that can be detected within the pixel, as a function of redshift (and of position on the CCD). For this data set, the size distribution of galaxies does not enter as a nuisance in the data analysis, instead it enters as an important element of any theoretical model, an element which directly affects the predicted function $h(x)$.  \\citet{l02} obtained the function $h(x)$ over a limited range of $x$ values in each of ten redshift bins. They then fit the curves to a broken power law model, allowing one of the parameters of the model to vary with redshift. The choice of a broken power law and of the allowed types of variation with redshift was motivated by the appearance of the data, not by any physical model of galaxy formation. Due to surface brightness dimming, at high redshift the observations can only detect the upper end of pixels, i.e., those with the highest SFR intensities. Indeed, at all redshifts $z>2$, most of the cosmic SFR occurs at $x$ values that are below the detection limit, and thus the total cosmic SFR is sensitive to an extrapolation which depends on the assumed shape of the function $h(x)$. In this paper, we confront the data with a semi-analytic model based on the theoretical understanding of hierarchical galaxy formation in a $\\Lambda$CDM cosmology. We examine whether the model can explain the overall trends in the data, and we use the predicted, physically-motivated shape of $h(x)$ to carry out a measurement of the cosmic SFR.  The basic theoretical framework in which the matter content of the universe is dominated by CDM has recently received a major confirmation from measurements of the cosmic microwave background \\citep{Boomerang,Maxima,Dasy}. Based primarily on these measurements, in this paper we assume a $\\Lambda$CDM cosmology with parameters $\\Omm$ = 0.3, $\\Omega_\\Lambda$ = 0.7, $\\Omega_b = 0.05$, $\\sigma_8 = 0.8$, $n=1$, and $h=0.7$, where $\\Omm$, $\\Omega_\\Lambda$, and $\\Omega_b$ are the total matter, vacuum, and baryon densities in units of the critical density, $\\sigma_8$ is the root-mean-square amplitude of mass fluctuations in spheres of radius $8\\ h^{-1}$ Mpc, and $n=1$ corresponds to a primordial scale-invariant power spectrum.  Throughout this paper we express results in physical units in $\\Lambda$CDM. Note that \\citet{l02} calculated results in an $\\Omm=1$ cosmology, and expressed most quantities in units of $h=1$. We convert their measurements into our units and cosmology, using the redshift distribution of their galaxy sample. Note that $x$ depends on redshift but not on the cosmological matter content, since it is derived from observations of surface brightness. However, $h(x)$ and the cosmic SFR do change, through the change in both the proper area and the comoving volume corresponding to a given solid angle. Using $h=0.7$ reduces $h(x)$ by a factor of 1.4, and the conversion to $\\Lambda$CDM causes an additional decline by a factor which equals 1.3 for the $z=0$--0.5 bin, and rises up to 1.8 for all $z>2$ bins.  In the following section we present the details of our theoretical model. Readers primarily interested in the comparison to the data may go directly to the summary of the basic setup in \\S3, the results in \\S 4 and the conclusions in \\S 5. ",
        "conclusions": "We have predicted the distribution of SFR intensity based on models of hierarchical galaxy formation. We have found that these models provide a natural explanation for the observed trend of the increase in the typical SFR intensity with redshift, an increase which occurs despite the decrease in the typical mass of galaxies. The observed data of \\citet{l02} are consistent with a constant efficiency of turning gas into stars (best-fit $\\eta=2.3\\%$) and a constant duty cycle (best-fit $\\zeta=17\\%$). Thus, the data are consistent with the standard picture of the cosmic SFR rising from $z=0$ to $z=2$ and not increasing much further at $z>2$.  However, the $h(x)$ data are limited as a probe of star formation efficiency, since the broad spin-parameter distribution, along with the distribution of disk orientations, smoothes $h(x)$ into an approximate power law with only a gradual break. This break is difficult to measure, since only a limited range of $x$ values can be detected, especially at high redshift. Thus, two-parameter $\\chi^2$ contours show a near-degeneracy between the two fitted parameters, and much more data would be required in order to fit even more detailed theoretical models of $h(x)$.  Luminosity functions of galaxies can be obtained for relatively large populations, using ground-based observations; also, if the total luminosity of each galaxy can be reliably measured, then this integrated quantity may allow a more robust comparison with theoretical models.  However, data in the form of $h(x)$ do have a number of advantages. The effects of surface-brightness dimming and of the size distribution of galaxies can both be directly incorporated in the analysis, unlike the case of galaxy luminosity functions where both effects enter the selection function in ways that are difficult to model. In addition, other corrections may also be simpler; e.g., extinction may well depend directly on $x$, and not on the total luminosity of a galaxy, though high resolution images in many wavebands would be needed in order to determine the level of extinction separately in every pixel. Indeed, the upcoming Next Generation Space Telescope should produce a great leap forward for this type of data. It should probe a much wider range of $x$ values, with large numbers of galaxies detected over redshifts up to 10 and beyond. It should also provide high angular resolution over a wide range of optical and infrared wavelengths, which will likely allow it to overcome the limitations of current data."
    },
    "0201/astro-ph0201347_arXiv.txt": {
        "abstract": "Using the imaging polarimeter for the Submillimeter Common User Bolometric Array at the James Clerk Maxwell Telescope, we have detected polarized thermal emission at 850 \\micron\\ from dust toward three star-forming core systems in the Orion B molecular cloud: NGC 2071, NGC 2024 and LBS 23N (HH 24).  The polarization patterns are not indicative of those expected for magnetic fields dominated by a single field direction, and all exhibit diminished polarization percentages toward the highest intensity peaks.  NGC 2024 has the most organized polarization pattern which is centered consistently along the length of a chain of 7 far-infrared sources.  We have modeled NGC 2024 using a helical field geometry threading a curved filament and also as a magnetic field swept up by the ionization front of the expanding \\hii\\ region.  In the latter case, the field is bent by the dense ridge, which accounts for both the polarization pattern and existing measurements of the line-of-sight field strength toward the northern cores FIR 1 to FIR 4.  The direction of the net magnetic field direction within NGC 2071 is perpendicular to the dominant outflow in that region. Despite evidence that line contamination exists in the 850 \\micron\\ continuum, the levels of polarization measured indicate that the polarized emission is dominated by dust. ",
        "introduction": "\\label{p3:intro}  The Orion B (L1630) molecular cloud, at a distance of 415 pc \\citep{ant82}, is one of the nearest giant molecular clouds and is an active site of low- to high-mass star formation.  It was one of the first clouds to be systematically studied for dense cores by \\citet{lbs91}, who found that massive star formation takes place only in the five largest clumps, which together make up more than 50\\% of the mass of dense gas.  We have chosen three of these five clumps, NGC 2071IR (LBS 8), NGC 2024 (LBS 33), and HH24 (LBS 23), for the current study.  A fourth region, NGC 2068, contains a string of substantially smaller, fainter cores connected by weak dusty filaments \\citep{mit01}.  Polarimetry of this region will be presented in a forthcoming paper \\citep[Paper IV in this series]{mw01}.  Although the three regions have comparable gas masses, ranging from 230-460 M$_\\odot$ \\citep{lbs91}, they have very different star formation properties.  NGC 2071IR lies four arcminutes north of the reflection nebula NGC 2071.  This extended submillimeter source consists of a cluster of at least nine embedded infrared stars \\citep{wal93} with a combined infrared luminosity of 520 L$_\\odot$ \\citep{but90}.  The source IRS 3 is thought to be the driving source of a massive bipolar molecular outflow \\citep{bal82,sne84,mor89,eis00}.  \\citet{hou00} infer alignment between the outflow and its magnetic field by comparison of spectral lines of neutral and ionic species.  Shocked molecular hydrogen \\citep{bla82} and H$_2$O masers \\citep{gen79} are also seen towards this region, which is in a later evolutionary stage than NGC 2024 \\citep{lau96}.  \\citet{eis00} has also documented several other outflows in the region.  Submillimeter maps of this region have been published by several authors \\citep{mit01,mot01,joh01}.  Based on a comparison between submillimeter continuum and CO(3-2) and HCO$^+$ line data, \\citet{mot01} suggest that $> 20$\\% of the 850 \\micron\\ emission in the outflow region could originate from line contamination.  By contrast with the other two regions, the HH 24-26 (LBS 23) cores have relatively little extended submillimeter emission and are mostly compact \\citep{mit01,joh01,lau96,lis99} and cold ($<$10K, \\citet{chi93}).  Of the twelve condensations identified by \\citet{lis99}, our polarimetric image covers LMZ 2, 3, and 4.  All three of these cores have 3.6 cm continuum sources \\citep{baw95,gib99}.  LMZ 3 (also known as HH24MMS, \\citet{chi93}) is a class 0 protostar \\citep{baw95}, while LMZ 4 is a T Tauri star with a known CO outflow \\citep{sne82,gib93}.  These are all indicators of low-mass star formation.  NGC 2024 has an associated \\hii\\ region and is the most prominent star formation region in Orion B, associated with a massive cluster, ionizing B stars, and stars at all phases of evolution \\citep{mez88,lad91,cc96}.  The submillimeter continuum emission was discovered by \\citet{mez88}.  The emission arises from a dense ridge of gas and dust behind the \\hii\\ region, determined from the velocity of associated gas \\citep{cru86,bar89}, and consists of at least seven sources aligned along a ridge, similar to OMC-3 in Orion A \\citep{joh99,chi97}.  Relatively more mass exists in the filamentary gas in the latter region.  Two of these cores (FIR 4 and FIR 5) are the origins of unipolar molecular outflows, one of which is very highly collimated and very extended \\citep{san85,ric92,cc96}, while the FIR 6 core exhibits a compact outflow \\citep{cc96} and contains a water maser \\citep{gen77}, a signature of intermediate-mass protostars.  The rest of the cores show no sign of star formation activity \\citep{vis98}.  \\citet{fuk00} present a numerical simulation of triggered star formation along a filament by the expansion of an \\hii\\ region and apply their analysis to NGC 2024 in particular.  In their picture, compression triggers two cores (i.e., FIR 4 and FIR 5) to collapse, and then at a later time, further collapse is triggered further up and down the filament.  This sequence was observed by \\citet{cc96} in their study of outflows from NGC 2024, where the dynamical ages of outflows from FIR 5, FIR 4 and FIR 6 are $1.4 \\times 10^4$ yr, $2.6 \\times 10^3$ yr, and 400-3300 yr.  Magnetic fields play a crucial role in the process of star formation, through the magnetic support of molecular clouds, dissipation of angular momentum in accretion disks, and the generation of jets and outflows (see \\citet{hei93} and references therein).  Polarized thermal emission at submillimeter wavelengths from aligned dust grains traces the direction of the magnetic field structure projected onto the plane of the sky \\citep{hil88}.  Absorption and scattering are usually negligible in the submillimeter.  This is in contrast to optical and near infra-red polarimetry, where both of these processes can cause contamination.  With the recent development of focal plane bolometer arrays equipped with polarimeters, sensitive imaging polarimetry in the submillimeter is now possible.  Polarimeters functioning at 100 \\micron\\ (aboard the Kuiper Airborne Observatory), 350 \\micron\\ (at the Caltech Submillimeter Observatory), and 850 \\micron\\ (at the James Clerk Maxwell Telescope) have detected polarized emission from dust toward many Galactic molecular clouds. These include the well-studied OMC-1 core in Orion \\citep{sch98,sch97,cop00,dot00}; Sagittarius A \\citep{ait00}; Sagittarius B2 \\citep{dow98} and recently OMC-3 in Orion (\\citet[hereafter Paper I]{mw00}; \\citet[hereafter Paper II]{mwf01}; Dowell 2001, in preparation).  Numerous protostellar and starless cores have also been mapped (i.e., \\citet{hol96,war00}).  Polarimetry provides information only on the plane-of-sky magnetic field orientation, but no direct information about the magnetic field strength, since the degree of polarization is also dependent on other factors such as grain shape, degree of alignment, and composition.  To obtain information about the magnetic field strength requires observation of Zeeman splitting of molecular or atomic spectral lines, which additionally provides information about the direction of the field along the line of sight.  In this paper, we present the first submillimeter polarimetry of the NGC 2071 core and the LBS 23N region.  Zeeman observations do not exist toward either of these two clouds.  Far-infrared polarimetry at 100 \\micron\\ \\citep{dot00} and Zeeman splitting observations of OH \\citep{cru83,kaz86,crtg99} exist for the NGC 2024 ridge of cores.  In $\\S$ \\ref{p3:obs}, we describe the 850 \\micron\\ observations and the data reduction.  The data are presented in $\\S$ \\ref{p3:sec:results} and the polarization patterns are interpreted in $\\S$ \\ref{p3:sec:theory}.  We summarize our results in $\\S$ \\ref{p3:summary}. ",
        "conclusions": "\\label{p3:summary}  The polarization patterns observed in Orion B are dominated by orderly structure on scales at least as large as the areas mapped (with the possible exception of the LBS 23N region).  This result can only be consistent with magnetic fields that are ordered on similar scales or larger.  The filamentary clouds, NGC 2024 and LBS 23N, exhibit polarization patterns that are inconsistent with purely poloidal magnetic fields threading the filaments, or with fields simply threaded transversely to them.  Curved field lines are necessary to model the polarization data in NGC 2024.  The fact that the polarization pattern is symmetric about the dense ridge of cores suggests that there is a correlation between the presence of dense gas and the exhibited polarization pattern.  The polarization systematically decreases with intensity for all three regions mapped.  This result is in agreement with Paper II, which notes depolarization toward the axis of the \\intfil\\ in OMC-3, as well as most submillimeter and millimeter observations of star-forming regions (see \\citet{wga00} for a review).  This ``polarization hole'' effect is a {\\em global} property of our maps, and not just a local phenomenon that occurs for a few pixels near the very brightest emission.  NGC 2071 is a massive core forming multiple protostars.  Its polarization pattern is ordered and qualitatively similar to that of OMC-1, which is an even more massive core in Orion A.  In OMC-1, \\citet{sch98} interpreted their polarization data as being due to an hourglass magnetic field, with the field lines pinched due to the collapse of the core.  This interpretation is consistent with the flattening observed along the inferred field lines in OMC-1.  On the other hand, NGC 2071 does not show any flattening, which is inconsistent with an interpretation of the magnetic field threading this core as a dynamically significant hourglass field.  We note that if the vectors of NGC 2071 were rotated to infer a net field direction, this direction would be {\\it perpendicular} to the most powerful outflow in the region.  Alternatively, our map could be interpreted as resulting from a field that is predominantly toroidal about the axis of symmetry.  Although CO $J=3-2$ emission can contaminate SCUBA 850 \\micron\\ data, we estimate it can contribute a maximum of 1\\% of polarized emission when optically thick.  Since much higher polarizations are measured, we conclude that the polarizations measured in NGC 2071 are dominated by dust emission.  Application of the Chandrasekhar \\& Fermi method to this core yields a mean field strength estimate of 56 $\\mu$G, and we find that half the magnetic energy is accounted for by the mean field component.  LBS 23N exhibits the least orderly polarization patterns of the three regions studied.  Most of the polarization vectors are aligned in a north-south orientation, particularly to the east of the cores.  The cores themselves are significantly depolarized.  Near the southern boundary of the map, the vectors rotate by $90^\\circ$ to an east-west orientation.  This abrupt change might be explained either by an extension of the Fiege \\& Pudritz model (see \\S \\ref{p3:sec:LBS23N}), or by a smaller filament orthogonal to the main filament.  More data to the south of LBS 23N would be needed to distinguish between these possibilities.  The NGC 2024 ridge of cores presents the most interesting polarization data of this paper.  The polarization is strong everywhere except the cores, which are significantly depolarized.  Our map agrees very well with the 100 \\micron\\ polarization data obtained with the KAO Stokes polarimeter \\citep{dot00}, even though these wavelengths likely probe different dust temperatures.  Interestingly, the region between FIR 4 and FIR 5 exhibits diminished intensity but no polarization is detected there.  We are able to successfully model the polarization pattern using a helical field geometry threading a curved filament. However, given the overall geometry of the region -- an \\hii\\ region expanding into the molecular cloud in which the cores are embedded -- we slightly prefer a model in which the ionization front has swept up the magnetic field and is now stretching it around the ridge of dense cores.  Since this model contains an unpolarized filament, it may be oversimplified and a more complete, complex picture will include both the polarized ionization front and a polarized filament.  We note that our analysis of the NGC 2024 polarization data have been aided immensely by the existence of Zeeman maps of the line-of-sight field strength toward part of this region published by \\citet{crtg99}. Unfortunately, very few regions have been successfully mapped using Zeeman splitting \\citep{cru99a}.  This is an area where future technological advances in instrumentation would greatly improve our understanding of magnetic fields in molecular clouds."
    },
    "0201/astro-ph0201037_arXiv.txt": {
        "abstract": " ",
        "introduction": "The atmospheric neutrino results reported by the  Super-Kamiokande \\cite{sk} and MACRO \\cite{mac,mac2k} experiments indicate that neutrinos oscillate, these data being consistent with $\\nu_{\\mu} \\leftrightarrow \\nu_{\\tau}$ oscillations. The small value of the difference of the squared masses ($5 \\times 10^{-4}{\\rm eV}^2  \\leq \\Delta m^{2} \\leq 6 \\times 10^{-3}{\\rm eV}^2$) and the strong mixing angle ($ \\sin^{2} 2\\theta \\geq 0.82$) suggest that these neutrinos are nearly equal in mass as predicted by many models of particle physics beyond the standard model. Also, the LSND experiment \\cite{Ath98} supports $\\nu_{\\mu}\\leftrightarrow \\nu_{e}$ oscillations ($\\Delta m^{2} \\leq 0.2$eV$^{2}$); some solar neutrino experiments \\cite{Bah98} suggest that $\\nu_e$ could oscillate to a sterile neutrino $\\nu_e \\leftrightarrow \\nu_s$ ($\\Delta m^{2} \\simeq 10^{-5}$eV$^2$). The direct implication of neutrino oscillations is the existence of a non-zero neutrino mass in the eV range or lower, and consequently a not negligible hot dark matter (HDM) contribution $\\Omega_{\\nu}~\\neq 0$ to  the total energy density of the universe. Recent works \\cite{Primack95,Gaw98,Gaw2000} show that the addition of a certain fraction of HDM component to the total energy density of the universe can lead to the agreement between the Cosmic Microwave Background (CMB) anisotropy power spectrum at the small scales and  the observations of the Large Scale Structure (LSS) of the universe as derived by the redshift galaxy surveys \\cite{Primack95}.  The CMB anisotropy pattern contains information related to the physical processes occurring before the last scattering of the CMB photons; the LSS data reflects the clustering regime effects in our local universe. These measurements can provide independent probes for the structure of the universe on similar comoving scales at different cosmological epochs. This paper discusses the implications of massive neutrinos properties for CMB and LSS measurements. ",
        "conclusions": "The standard Cold Dark Matter model has difficulties in matching the CMB and LSS measurements. Cosmological models involving a mixture of CDM and HDM particles, the CHDM models, are able to fit, the  excess large scale power seen in galaxy surveys and the CMB temperature fluctuations.  The results from neutrino oscillation experiments indicate a non-zero neutrino mass in the eV range, or lower. This  implies a non-negligible contribution of neutrinos to the total energy density of the universe. The experimental evidence indicating a present low matter density universe,  dominated by the vacuum energy ($\\Omega_m \\simeq 0.3$, $\\Omega_{\\Lambda}\\simeq 0.7$) and a higher Hubble parameter value ($H_0 \\simeq 65$ km s$^{-1}$Mpc$^{-1}$) are in agreement with cosmological models involving neutrinos if one considers the lepton asymmetry of the neutrino background.  At the times when the anisotropies were generated, neutrinos had significant interactions with the photons, baryons and cold dark matter particles only via gravity. Neutrinos can not cluster via gravitational effects on scales $k$ below the free streaming scale $k<k_{fs}$ (large wavelengths). For $k>k_{fs}$ the growth of the density perturbations is suppressed, the magnitude of this suppression depending on $\\Omega_m$, $\\Omega_{\\nu}$, $\\Delta m^2$ and $\\Delta \\xi_{\\nu}$.  The neutrino homogeneous quantities (density and pressure) and inhomogeneous quantities (density and pressure perturbations, the shear stress and the energy flux) are also changed in the presence of non-degenerated neutrino mixing, leaving imprints on the CMB temperature fluctuations and the matter density fluctuations power spectra.  The determination of the fundamental cosmological parameters with high precision CMB and LSS surveys requires complementary knowledge of neutrino properties from cosmic rays and long base-line experiments.   \\vspace{0.6cm} {\\bf Acknowledgements:} I acknowledge  the organizers and  V.~Berezinsky for useful discussions during this workshop."
    },
    "0201/astro-ph0201127_arXiv.txt": {
        "abstract": "We combine complementary datasets to constrain dark energy. Using standard Big Bang Nucleosynthesis and the observed abundances of primordial nuclides to put constraints on $\\Omega_Q$ at temperatures near $T \\sim 1MeV$, we find the strong constraint $\\Omega_Q(\\mbox{MeV}) < 0.045$ at $2\\sigma$ c.l.. Under the assumption of flatness, using results from Cosmic Microwave Background (CMB) anisotropy measurements, high redshift supernovae (SN-Ia) observations and data from local cluster abundances we put a new constraint on the equation of state parameter $w_Q < -0.85$ at $68 \\%$ c.l.. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0201/astro-ph0201311_arXiv.txt": {
        "abstract": "We have observed the J=5-4 ground state transition of FeO at a frequency of 153 GHz towards  a selection of galactic sources. Towards  the galactic center source SgrB2,  we see weak  absorption at approximately the velocity of other features towards this source ( 62 \\kms \\ LSR). Towards other sources, the results were negative as they were also for MgOH(3-2) and FeC(6-5). We tentatively conclude that the absorption seen toward SgrB2 is due to FeO in the hot ($\\sim $ 500~K) relatively low density absorbing gas known to be present in this line of sight. This is the first (albeit tentative) detection of FeO or any iron--containing molecule in the interstellar gas. Assuming the observed absorption to be due to FeO, we estimate [FeO]/[SiO] to be of order or less than 0.002 and [FeO]/[\\MOLH] of order $3\\, 10^{-11}$.  This is compatible with our negative results in other sources. Our results suggest that the iron liberated from grains in the shocks associated with  SgrB2 remains atomic and is not processed into molecular form. ",
        "introduction": "The failure to discover iron--bearing molecules in the interstellar medium is a long standing puzzle.  It is related to the general problem of the depletion of refractory elements (see e.g. Jenkins 1989, Weingartner and Draine 1999, Walmsley et al. 1999, Walmsley 2000) within both diffuse and dense molecular clouds.  There is good evidence that the degree of depletion is correlated with density in diffuse clouds and that it is extremely high within molecular clouds.  In fact, the abundance of gas phase silicon appears to be six orders of magnitude below the solar Si abundance in some circumstances (see also Ziurys, Friberg, \\& \\ Irvine  1989).  Moreover, Turner(1991) has put limits on the abundance of a variety of molecules containing Na, Si, Mg, Fe, and P showing that the case of silicon is not unusual.  The refractory elements  thought to be the main constituents of ``silicate grains'' are even more underabundant in the gas phase of dense molecular clouds than they are in the diffuse medium sampled by UV observations.  Nevertheless, silicon (essentially in the form SiO) is known to be present at a low level in  some molecular clouds associated with outflows (see Bachiller and Tafalla 1999, Bachiller et al. 2001, Codella et al. 2001) as well as in photon dominated regions (PDRs, Schilke et al. 2001).  The general interpretation of this is that a small fraction of Si is returned to the gas phase in shocks associated with star formation in molecular clouds (see e.g. Caselli, Hartquist, \\& \\ Havnes 1997, Schilke, Pineau des For\\^{e}ts \\& \\ Walmsley 1997). One might therefore naively expect iron and magnesium to be also liberated in such shocks and hence to be present in the molecular gas phase at the same low level. With this in mind, we have started an observational program searching for Mg/Fe containing species  associated with shocks.  This article describes a small search carried out with the IRAM 30-m telescope  for iron and magnesium containing species. In the course of this, we detected evidence for the presence of FeO in the molecular clouds seen in absorption along the line of sight towards the continuum source in the vicinity of the galactic center SgrB2-M. SiO is well known along this line of sight (Greaves, Ohishi, \\& \\ Nyman  1996, H\\\"{u}ttemeister et al. 1995, Peng, Vogel, \\& \\ Carlstrom 1995) and indeed is more  abundant in galactic center clouds in general (H\\\"{u}ttemeister et al. 1998) than in molecular clouds in the solar vicinity. The observed large line widths suggest that this may be due to shocks (e.g. Flower, Pineau des For\\^{e}ts \\& \\ Walmsley 1995) caused by cloud--cloud collisions due to shearing motions in the inner Galaxy. It is also possible that  the chemistry is affected by the hard X--ray sources in the vicinity of SgrB2  which heat and ionize the neighbouring molecular clouds (see Martin--Pintado et al. 2000). In any case, one expects a correlation between silicon and iron and hence it is reasonable to expect to find traces of FeO in such regions. In this letter, we  present the evidence that FeO has been detected in the interstellar medium and briefly mention some consequences for the chemistry of iron in molecular clouds. ",
        "conclusions": "We believe that we have detected FeO in absorption towards SgrB2. However, the feature which we have detected is weak and more confirmatory measurements are needed. Searches for other low excitation transitions of iron--bearing species would be useful.  A more detailed study of the interstellar chemistry of iron is also needed.  Even if our identification in SgrB2 turns out to be incorrect, one can conclude that in the SgrB2 absorbing layer as well as in the post-shock gas which one observes towards L1157, one has [FeO]/[SiO] less than 0.01. SiO is thought to be a major form of gas phase silicon and thus the SiO abundance gives a measure of silicon depletion. For FeO, the preliminary model calculations mentioned earlier suggest that iron is indeed produced by erosion in shocks but remains atomic in the post--shock medium.  Indeed [FeII] emission is well known in the shocks associated with the Orion outflow (Tedds, Brand, \\& \\ Burton 1999) and so this is quite plausible. We conclude therefore that erosion of silicate grains in high velocity (40\\kms ) shocks is a plausible explanation of our observation towards SgrB2."
    },
    "0201/astro-ph0201257_arXiv.txt": {
        "abstract": "Contrary to earlier expectations, several cosmic ray events with energies above $10^{20}$ eV have been reported by a number of ultra-high energy cosmic ray observatories. According to the AGASA experiment, the flux of such events is well above the predicted Greisen-Zatsepin-Kuzmin cutoff due to the pion production of extragalactic cosmic ray protons off the cosmic microwave background. In addition to the relatively high flux of events, the isotropic distribution of arrival directions and an indication of small scale clustering strongly challenge all models proposed to resolve this puzzle.  We discuss how the GZK cutoff is modified by the local distribution of galaxies and how astrophysical proton sources with soft injection spectra are ruled out by AGASA data. Sources with hard injection spectrum are barely allowed by the observed spectrum. If the most recent claims by AGASA that the highest energy events are due to clustered nuclei are confirmed, the most plausible explanation are astrophysical sources with very hard spectra such as extragalactic unipolar inductors. In addition, extragalactic magnetic fields need to be well below the current nano-Gauss upper limits. Alternatively, if the primaries are not nuclei, the need for new physics explanations is paramount.  We present an overview of the theoretical proposals along with their most general signatures to be tested by upcoming experiments. \\vspace{1pc} ",
        "introduction": "The future of ultra high energy cosmic ray physics looks extremely promising. The present state of observations is particularly puzzling and the necessary experiments to resolve these puzzles will be operating in the very near future. The puzzles begin with the lack of the predicted  Greisen-Zatsepin-Kuzmin (GZK) cutoff \\cite{g66,zk66}. Contrary to earlier expectations, cosmic rays with energies above $10^{20}$ eV have been detected by a number of experiments (for a review see  \\cite{nw00} and for a more recent update see \\cite{icrc01}).  If these particles are protons, they  are likely to originate in extragalactic sources, since at these high energies the Galactic magnetic field cannot confine protons in the Galaxy. However, extragalactic protons with energies above  a few times $10^{19}$ eV can produce pions through interactions with the cosmic microwave  background (CMB) and consequently lose significant amounts of energy as they traverse intergalactic distances.  Thus, in addition to the extraordinary energy requirements for astrophysical sources to accelerate protons to $\\ga \\ 10^{20}$ eV, the photopion threshold reaction  suppresses the observable flux above $\\sim 10^{20}$ eV. These conditions were expected to cause a natural high-energy limit to the cosmic ray spectrum known as the GZK cutoff \\cite{g66,zk66}.  As shown by the most recent compilation of the AGASA data \\cite{agasa01a}, the spectrum of cosmic rays does not end at the expected GZK cutoff. The significant flux observed above $10^{20}$ eV together with a nearly isotropic distribution of event arrival directions \\cite{agasa01b} challenges astrophysically based explanations as well as new physics alternatives  (see \\cite{bs00,o00}  and references therein). In addition, the reported small scale clustering \\cite{agasa01b} tends to rule out most scenarios.  This challenging state of affairs is stimulating both for theoretical investigations as well as experimental efforts. The explanation may hide in the experimental arena such as an over estimate of the flux at the highest energies. This explanation has been proposed by the HiRes collaboration based on an analysis of their monocular data \\cite{hires01}. Even if this were the case, events above the GZK cutoff are also observed by HiRes. The Mono HiRes data looks more like the GZK {\\it feature}, as discussed in the next section, followed by indications of new sources at energies above the feature. Events past $10^{20}$ eV pose theoretical challenges which will be explained in the future by either astrophysically novel sources or new fundamental physics. ",
        "conclusions": "Next generation experiments such as the Pierre Auger Project which is now under construction,  the proposed  Telescope Array, and the EUSO project and the OWL  satellites  will significantly improve the data at the extremely-high end of the cosmic ray spectrum.  With these observatories a clear determination of the spectrum and spatial distribution of UHECR sources is within reach. In addition, the observations of UHE neutrinos in horizontal showers promises to open a new window into the workings of our Universe.   The lack of a GZK cutoff should become clear with Auger and most extragalactic Zevatrons may be ruled out. The observed spectrum will distinguish Zevatrons from new physics models by testing the hardness of the spectrum and the effect of propagation.  Fig. 8 shows how clearly Auger will test the spectrum independent of their clustering properties. The cosmography of sources should also become clear and able to discriminate  between plausible populations for UHECR sources. The  correlation of arrival directions  for events with energies above $10^{20}$ eV  with some known structure such as the Galaxy, the Galactic halo, the Local Group or the Local Supercluster would be key in differentiating between different models. For instance, a correlation with the Galactic center and  disk should become apparent at extremely high energies for the case of young neutron star winds, while a correlation with the large scale galaxy distribution should become clear for the case of quasar remnants. If SHRs  are responsible for UHECR production, the arrival directions should correlate with the dark matter distribution and show the halo asymmetry. For these signatures to be tested, full sky coverage is essential. Finally,  an excellent discriminator would be an unambiguous composition determination  of the primaries. In general, Galactic disk models  invoke iron nuclei to be consistent with the isotropic distribution, extragalactic Zevatrons tend to favor proton primaries, while photon primaries are more common for early universe relics.  The hybrid detector of the Auger Project should help settle the present disparity between HiRes and AGASA by cross calibrating the two techniques. It will also determine the composition by measuring the depth of shower maximum and the ground footprint of the same shower. AGASA seems to detect a hint of composition shifts at the highest energies. This would be quite a surprising development.  In sum, the future looks very promising. The solution to the UHECR mystery as well as the birth of UHE neutrino astronomy is coming with the next generation of experiments which are under construction such as Auger or in the planning stages such as the Telescope Array, EUSO, and OWL."
    },
    "0201/astro-ph0201531_arXiv.txt": {
        "abstract": "We present a model for magnetic structured coronae above accretion discs. On the shortest timescales, spatially and temporally correlated coronal flares can explain X-ray temporal and spectral variability observed in Seyfert galaxies. In particular, power density spectra, flux-spectral index and flux-variance correlations are naturally accounted for by the model. More dramatic spectral variations (i.e. state transitions in GBHC) are associated with parameters varying on longer timescales, such as accretion rate, coronal strength  or geometry of the inner disc. In the framework of the standard Shakura--Sunyaev accretion disc theory, here we discuss why energetically dominant coronae at low accretion rates are ideal sites for launching powerful MHD driven outflows. Then, if the outflow is radiatively inefficient, then so is the source overall, even without advection being relevant  for the dynamics of the accretion flow. This could be an alternative scenario for LLAGN and GBHC in their low/hard state, and may have consequences for our understanding of the accretion history of the universe. ",
        "introduction": "The hard X-ray spectra, the properties of the reflection features and of rapid time variability in accreting black holes can all be considered as indications of the presence of a structured, hot, optically thin component in the inner part of the accretion flow (the {\\it corona}), situated above a cold, geometrically thin, disc. Many models of the accretion disc corona have been proposed in recent years, which are able to fit the observed {\\it time averaged} spectra in terms of Comptonization of soft photons in the hot corona. However, many uncertainties regarding the actual geometry of the inner accretion flow remain unsolved (see Done 2001 for a review). Here we present a model for a structured, magnetic corona above a standard, geometrically thin and optically thick disc, that can satisfy observational constraints on both long and short time variability of these systems.  Let us briefly sketch the overall energetics of an accretion disc-corona system. The total power released by the accreting gas is defined as $L\\equiv \\dot m L_{\\rm Edd}=4\\pi G M m_{\\rm p}\\dot m c/\\sigma_{\\rm T}$. To be fairly general, we assume that the coronal power generated by the disc (which is a fraction $f$ of the total: $L_{\\rm c}=f\\dot m L_{\\rm Edd}$) can be either dissipated locally to heat the corona, and ultimately radiated away as hard X-rays with a luminosity $L_{\\rm H}=(1-\\eta)L_{\\rm c}$, or used to launch an outflow with power $L_{\\rm j}=\\eta L_{\\rm c}$. The main characteristic of our  model are the following: in the hot corona the energy is stored in a strong, highly intermittent magnetic field (amplified in the turbulent disc and buoyantly expelled in the vertical direction). Magnetic energy dissipation occurs at the smallest end of the turbulent-energy cascade. Such small flares heat the corona (with a power $L_{\\rm H}$) and can trigger an avalanche in their immediate neighborhood (\\cite{amerloni-C2:pf99}), creating a bigger active region, and producing the observed flares in the lightcurves. In Sec.~\\ref{amerloni-C2_sec_short}, a stochastic model (the so-called thundercloud model, Merloni \\& Fabian, 2001, MF1) of the short-time variability is discussed.  On longer timescales, variations of the fraction of power released into the corona, $f$, and of the fraction of coronal power used to launch an outflow, $\\eta$, may be associated, for example, to changes of the accretion rate and/or of the geometry of the inner disc. In Sec.~\\ref{amerloni-C2_sec_long} we show that, under reasonable assumptions on the nature of the disc viscosity, the strength of the corona $f$ increases as the accretion rate decreases. Also we discuss reasons why an energetically dominant corona can be the site where powerful outflows are produced. ",
        "conclusions": "We have presented a model to explain spectral and temporal variability on the smallest timescales in the X-ray emission from Seyfert Galaxies and GBHC. We have simulated X-ray light-curves that reproduce the observed PDS properties and the spectral variability. The basic geometric properties of the corona we propose are the following: \\begin{itemize} \\item{The corona must not be uniform, but structured and heated intermittently (flares);} \\item{The spatial and temporal distribution of the flares are not random, but proceed in correlated trains of events (avalanches);} \\item{The size of the avalanches determines the size of the active regions, which are distributed as a power-law; larger avalanches are more luminous and have softer spectra.} \\end{itemize}  On longer timescales, the evolution of the corona is governed by the evolution of the accretion rate and/or of the inner disc geometry. At low accretion rate, the strength of a magnetic corona produced by buoyant magnetic flux tube amplified in an underlying standard accretion disc increases. If the energy in the corona, as we suggest, is stored in the magnetic field, and the height of a reconnection site is much larger than its size, which is of the order of the disc thickness, powerful MHD outflows can be launched from the inner corona."
    },
    "0201/astro-ph0201477_arXiv.txt": {
        "abstract": "We study the dust depletion pattern in eight well separated components of the $z_{\\rm abs}$~=~1.973, log~$N$(H~{\\sc i})~=~20.83, damped Lyman-$\\alpha$ system toward Q~0013$-$004, four of which have detectable H$_2$ absorption. The apparent correlation between the abundance ratios [Fe/S] and [Si/S] in the components indicates that the abundance pattern is indeed due to dust-depletion. In particular, we find evidence for depletion similar to what is observed in cold gas of the Galactic disk ([Fe/Zn]~=~$-$1.59, [Fe/S]~=~$-$1.74, [Zn/S]~=~$-$0.15, [Si/S]~=~$-$0.85) in one of the weakest components  in which molecular hydrogen is detected with log~$N$(H$_2$)~$\\sim$~16.5. This is the first time such a large depletion is seen in a DLA system.  Extinction due to this component is negligible owing to small total \\hi column density, log~$N$(\\hi)~$\\leq$~19.4. This observation supports the possibility that current samples of DLA systems might be biased against the presence of cold and dusty gas along the line of sight. \\par\\noindent The overall metallicities of this peculiar DLA system in which O~{\\sc i} and C~{\\sc ii} are spread over $\\sim$1050~km~s$^{-1}$ are [P/H]~=~$-$0.64, [Zn/H]~=~$-$0.74 and [S/H]~=~$-$0.82 relative to solar. The sub-DLA system at $z_{\\rm abs}$~=~1.96753 has [P/H]~$>$~0.06, [Zn/H]~$>$~$-$0.02 and [S/H]~$>$~$-$0.18. The overall molecular fraction is in the range $-$2.7~$<$~log~$f$~$<$~$-$0.6, which is the highest value found for DLA systems. H$_2$ is detected in four components at $-$615, $-$480, 0 and 85~km~s$^{-1}$ relative to the strongest component at $z_{\\rm abs}$~=~1.97296. CO is not detected (log~$N$(CO)/$N$(H~{\\sc i})~$<$~$-$8) and HD could be present at $z_{\\rm abs}$~=~1.97380. \\par\\noindent We show that the presence of \\h2 is closely related to the physical conditions of the gas: high particle density together with low temperature. The observed excitation of high $J$ H$_2$ levels and the molecular fraction show large variations from one component to the other suggesting that the UV radiation field is highly inhomogeneous throughout the system. Gas pressure, estimated from C~{\\sc i} absorptions, is larger than what is observed in the ISM of our Galaxy. This, together with the complex kinematics, suggests that part of the gas is subject to high compression due to either collapse, merging and/or supernovae explosions. This is probably a consequence of intense star-formation activity in the vicinity of the absorbing gas. ",
        "introduction": "The amount of dust present at high redshift has important consequences on the physics of the gas. In addition, dust directly affects our view of the high redshift universe through extinction. Therefore, the presence of dust in damped Lyman-$\\alpha$ (hereafter DLA) systems, that contain most of the neutral hydrogen in the universe, can have significant consequences. Although the presence of dust in DLA systems has been claimed very early (Pei et al. 1991), the issue has remained controversial. Indeed, Lu et al. (1996) have questioned the idea that the overabundance of Zn compared to Cr or Fe observed in DLA systems (e.g. Pettini et al. 1997) is due to selective depletion onto dust-grains and have argued that the overall abundance pattern observed in DLA systems is indicative of Type II supernovae enrichment instead. In recent years several studies have shown that both effects, dust-depletion and peculiar nucleosynthesis history, should be invoked to explain the abundance pattern (Vladilo 1998, Prochaska \\& Wolfe 1999, Ledoux et al. 2001a). However, the lack of statistics and the wide variety of objects that can give rise to DLA systems, namely dwarf galaxies (Centuri\\'on et al. 2000), large disks (Prochaska \\& Wolfe 1997, Hou et al. 2001), galactic building blobs (Haehnelt et al. 1998, Ledoux et al. 1998) etc., with, for each of these objects, its own history, prevent us to have a clear picture of the nature of DLA systems. Nevertheless, all studies conclude that the dust content of DLA systems is small. However, it is possible that the current sample of DLA systems is biased against high-metallicity and dusty systems. Indeed, Boiss\\'e et al. (1998) have noticed that there is a lack of systems with large $N$(H~{\\sc i}) and large metallicity. Very recent investigation of an homogeneous sample of radio-selected quasars shows that the dust-induced bias cannot lead to underestimate the H~{\\sc i} mass in DLA systems by a large factor (Ellison et al. 2001). However even a factor of two could change our understanding of DLA systems. \\par\\noindent An obvious way to search for DLA systems with large amount of dust is to select those where molecules are detected as these molecules form predominantly at the surface of dust grains. However, it may not be so simple as it has been shown that the presence of H$_2$ is not only related to the dust-to-metal ratio but is mostly dependent on the physical conditions of the gas. First of all, H$_2$ is detected when the particle density is large (Petitjean et al. 2000, Ledoux et al. 2001b). In any case, the system at $z_{\\rm abs}$~=~1.973 toward Q~0013$-$004 is a good target as very strong molecular absorption lines have been identified by Ge \\& Bechtold (1997, see also Ge et al. 1997). ",
        "conclusions": "The DLA system at $z_{\\rm abs}$~=~1.973 toward Q~0013$-$004 is peculiar in several aspects. Absorption lines from metal species are spread over about 1000~km~s$^{-1}$; in particular O~{\\sc i} and C~{\\sc ii} span 1050~km~s$^{-1}$. The velocity structure indicates the presence of two main sub-systems centered at $z$~$\\sim$~1.973 and 1.9674, separated by $\\sim$~550~km~s$^{-1}$ with, respectively, log~$N$(H~{\\sc i})~=~20.83 and $<$19.4 and [S/H]~=~$-$0.81 and $>-$0.18. \\par\\noindent The low-ionization gas is conspicuous in the system. There is clear evidence that all species are at most twice ionized in the $z_{\\rm abs}$~=~1.96822 component. This means that if photo-ionization dominates, there is probably very few photons with energy larger than 20~eV. More generally, $N$(X$^{+}$)/$N$(X$^{2+}$)~$>$~1 over the entire $z$~$\\sim$~1.973 system. This ionization state could reveal gas ionized by slow shocks. This idea is reinforced by the high pressure measured in a few components from C~{\\sc i} absorptions. \\par\\noindent \\h2 is detected in four main components, two very strong components (log~$N$(\\h2)~$>$~17) in the $z$~$\\sim$~1.973 system and two weaker components (log~$N$(\\h2)~$\\sim$~16) in the $z$~$\\sim$~1.9674 system. The total column density is 17.8~$<$~log~$N$(\\h2)~$<$~20.0 and therefore the mean molecular fraction, $f$~=~2$N$(\\h2)/(2$N$(\\h2)~+~$N$(H~{\\sc i})), is in the range $-$2.7~$<$~log~$f$~$<$~$-$0.6 which is the highest molecular fraction observed in DLA systems. \\par\\noindent The analysis of the $N$(C~{\\sc i}$^*$)/$N$(C~{\\sc i}) ratio in different components indicates that whenever \\h2 is detected, the particle density is high ($n_{\\rm H}$~$>$~30~cm$^{-3}$). High density is also found for components without any H$_2$ absorption detected. There is a hint for the depletion of metals within the components to be correlated to the $N$(C~{\\sc i}$^*$)/$N$(C~{\\sc i}) ratio. This suggests that depletion onto dust-grains could be larger for denser gas. \\par\\noindent The component at $z_{\\rm abs}$~=~1.96822 shows evidence for large depletion of iron and silicon relative to sulfur and zinc ([Fe/Zn]~=~$-$1.59, [Fe/S]~=~$-$1.74, [Zn/S]~=~$-$0.15, [Si/S]~=~$-$0.85) similar to what is observed in cold gas of the Galactic disk. The corresponding dust-extinction is small in this case because, although H$_2$ is detected with log~$N$(H$_2$)~$\\sim$~16.5, the H~{\\sc i} column density is small, log~$N$(H~{\\sc i})~$<$~19.4, in the component. This is direct evidence for a considerable fraction of heavy elements being locked into dust-grains, and, as a consequence, this supports the idea that the current sample of DLA systems might be biased against the presence of cold and dusty gas along the line of sight. Note that the rest of the gas shows a depletion pattern close to that of warm gas in the Galactic halo. \\par\\noindent The overall kinematics of the system with the two strong sub-systems separated by 550~km~s$^{-1}$ suggests that the line of sight is passing through one or several objects in strong interaction. The velocity structure of the subsystem at $z_{\\rm abs}$~=~1.9676 with the presence of molecular gas at $+100$ and $-150$~km~s$^{-1}$ from the center of the high-ionization absorptions strongly suggests the expansion of a spherical shell. All this, together with the strong inhomogeneity of the UV field, the high pressure in the gas and the high metallicities strongly suggests that intense star-formation activity is occuring in the vicinity of the system which should be revealed by deep imaging in the field."
    },
    "0201/astro-ph0201194_arXiv.txt": {
        "abstract": "We report the serendipitous detection in high--resolution optical spectroscopy of a strong, asymmetric \\lya\\ emission line at $z=5.190$. The detection was made in a 2.25 hour exposure with the Echelle Spectrograph and Imager on the Keck II telescope through a spectroscopic slit of dimensions 1\\arcsec\\ $\\times$ 20\\arcsec. The progenitor of the emission line, J123649.2$+$621539 (hereafter ES1), lies in the Hubble Deep Field North West Flanking Field where it appears faint and compact, subtending just 0\\farcs3 (FWHM) with $I_{\\mbox{\\tiny AB}} = 25.4$.  The ES1 \\lya\\ line flux of $3.0 \\times 10^{-17}$ ergs cm$^{-2}$ s$^{-1}$ corresponds to a luminosity of $9.0 \\times 10^{42}$ ergs s$^{-1}$, and the line profile shows the sharp blue cut--off and broad red wing commonly observed in star--forming systems and expected for radiative transfer in an expanding envelope. We find that the \\lya\\ profile is consistent with a galaxy--scale outflow with a velocity of $v > 300$ km s$^{-1}$.  This value is consistent with wind speeds observed in powerful local starbursts (typically $10^2$ to $10^3$ km s$^{-1}$), and compares favorably to simulations of the late--stage evolution of \\lya\\ emission in star--forming systems. We discuss the implications of this high--redshift galactic wind for the early history of the evolution of galaxies and the intergalactic medium, and for the origin of the UV background at $z > 3$. ",
        "introduction": "\\label{intro}  Following the epoch of recombination, the Universe settled into the comparatively dormant dark ages, during which the primordial glow had begun to fade but the first present--day astronomical objects had yet to form.  This tranquil epoch proved short--lived, however, as the formation of the first stars and quasars ushered in the first era of cosmological heating and enrichment at $z < 20$ \\citep[e.g.][]{gnedin97, haiman97, haiman98, ostriker96, valageas99}.  Evidence of these processes in the form of galaxy--scale outflows is abundant in spectroscopy of the high--redshift Universe.  Both optical/IR spectra of the $z \\sim 3$ Lyman--break population \\citep[e.g.][]{pettini01} and optical spectra of lensed \\lya--emitting galaxies at $z > 4$ \\citep{frye01} show metal absorption lines which are blueshifted by hundreds of km s$^{-1}$ with respect to the stellar rest frame of the galaxy, and \\lya\\ emission lines which are shifted similarly redward. These observations, as well as the characteristic P--Cygni profile of the \\lya\\ emission lines (e.g.\\ Bunker, Moustakas, \\& Davis 2000; Dey et al.\\ 1997, 1998; Dickinson 1998; Ellis et al.\\ 2001, Lowenthal et al.\\ 1997; Weymann et al.\\ 1998) paint a coherent picture of optically thick expanding regions surrounding star--forming galaxies, most naturally driven by the starbursts that render them visible in the first place \\citep[e.g.][and references therein]{heckman00}.  We present a direct observation of such an outflow at $z=5.190$ in high--resolution optical spectroscopy of the serendipitously detected star--forming galaxy J123649.2$+$621539 (hereafter ES1, for Echelle Spectrograph and Imager serendipitous detection number one). The sharp blue cut--off and broad red wing of the ES1 \\lya\\ emission line are consistent with the profile expected for the transfer of line radiation in an expanding envelope \\citep[e.g.][]{surdej79}.  The suggested outflow velocity of $v > 300$ km s$^{-1}$ is in broad agreement with simulations of the late evolution of \\lya\\ emission and absorption in star--forming galaxies \\citep[e.g.][]{tenorio99}, and is consistent with observations of powerful nearby starbursts \\citep*[e.g.][]{heckman90}. The spectrum of ES1 therefore presents evidence for both a high star--formation rate and a high--redshift, starburst--driven galactic wind, fitting well within the expectations of current models for the early history of galaxy formation.  In \\S \\ref{observation} we discuss the detection of ES1 and we give a description of the spectrum and the available archival imaging.  In \\S \\ref{line} we detail the properties of the ES1 \\lya\\ emission line, and we present a model for the emission line profile consistent with the expanding shell scenario introduced above. We conclude in \\S \\ref{discussion} with a discussion of the implications of the evidence for a strong outflow in ES1 for both the evolution of galaxies and the intergalactic medium (IGM)  at high redshift, and for the origin of the UV background at $z > 3$.  Throughout this paper we adopt the currently favored $\\Lambda$--cosmology of $\\Omega_{\\mbox{\\tiny M}} = 0.35$ and $\\Omega_\\Lambda = 0.65$, with $H_0 = 65$ km s$^{-1}$ Mpc$^{-1}$ \\citep[e.g.][]{riess01}.  At $z=5.190$, such a universe is only 1.10 Gyr old --- corresponding to a look--back time of 92.1\\% of the age of the Universe --- and an angular size of 1\\farcs0 corresponds to 6.31 kpc. ",
        "conclusions": "\\label{discussion}  The spectral profile of the \\lya\\ emission of ES1 presents evidence for a galaxy--scale outflow with a velocity of $v > 300$ km s$^{-1}$.  Of course, as \\citet{heckman00} caution, the outflow rate of a galactic wind cannot necessarily be interpreted as the rate at which mass or energy {\\it escapes} into the IGM, since the observable manifestation of an outflow may be produced by material still deep inside the gravitational potential well of the galaxy halo.  Nonetheless, the outflow velocity estimated for ES1 far exceeds the escape speed of a nominally low mass ($M < 10^{10} M_\\sun$) pregalactic fragment, consistent with the general observation that hot gas can readily escape from dwarf galaxies, though perhaps not from more massive systems \\citep{heckman00,heckman00p,martin99}.  This conclusion bears on a host of cosmological issues surrounding the evolution of galaxies and the IGM at high redshift.  Foremost, it suggests that processed material from ES1 will become available to the IGM, potentially providing the enrichment necessary to account for the amount of metals there observed.  Indeed, recent observations of \\ion{C}{4} absorption systems along the lines--of--sight to lensed QSOs call for enrichment at increasingly high redshift, beyond even $z > 5$ (e.g.\\ Aguirre et al.\\ 2001; Rauch, Sargent, \\& Barlow 2001). Additionally, both detailed observations and careful theoretical studies demand a mechanism for pre--heating the material out of which galaxy clusters ultimately collapse and become bound \\citep[e.g.][and references therein]{kaiser91,mushotzky97}. Here again, galaxy--scale outflows at high redshift are the likely culprit \\cite[e.g.][]{renzini93}.  Finally, galactic winds have proved important in reproducing the faint--end slope of the observed field galaxy luminosity function in semi--analytic models of galaxy formation. Outflows are invoked to suppress star--formation in low--mass dark matter halos, either via the direct escape of gas--phase baryons in the outflow itself \\citep[e.g.][]{somerville99}, or by ram pressure stripping of the gas--phase baryons by energetic winds from neighboring galaxies \\citep{scannapieco01}.  As a somewhat speculative conclusion, we now consider the expected correlation between strong galactic outflows and the escape of Lyman continuum radiation from star--forming galaxies.  This correlation bears directly on the much--debated physical nature and relative contributions of the sources which comprise the UV background, as a significant contribution by sources other than QSOs is required at high redshift, owing to the rapid decline in the space density of optical and radio--loud quasars at $z > 3$ \\citep*{bianchi01,madau99}.  It is likely that star--forming galaxies fill this niche.  From the theoretical standpoint, mechanical energy deposition in the form of supernovae and stellar winds is expected to result in an over--pressured cavity of hot gas inside star--forming galaxies.  In galaxies for which the star--formation rate per unit area $\\Sigma_\\ast \\geq 10^{-1} M_\\sun \\; \\mbox{yr}^{-1} \\; \\mbox{kpc}^{-2}$, the superbubble ultimately expands and bursts out into the galaxy halo, allowing for the escape of hot gas and facilitating the leak of Lyman continuum photons \\citep{heckman00,tenorio99}.  Of course, as the superbubble will expand in the direction of the vertical pressure gradient, the burst is expected to take the form of a weakly collimated, bipolar wind.  Hence, the leak of UV radiation may depend sensitively on not only the distribution of neutral gas and dust in the galaxy interstellar medium, but on the inclination of the system. Nonetheless, from the observational standpoint, \\citet*[][hereafter SPA01]{steidel01} report the detection of significant Lyman continuum emission in a composite spectrum of 29 Lyman--break galaxies at $\\langle z \\rangle = 3.40 \\pm 0.09$, suggesting an escape fraction\\footnote{Here, \\fesc\\ is the fraction of emitted 900 \\AA\\ photons that escapes the galaxy without being absorbed by interstellar material, normalized by the fraction of emitted 1500 \\AA\\ photons which similarly escapes.  As SPA01 point out, this definition differs from definitions encountered elsewhere, which typically consider only the fraction of emitted 900 \\AA\\ photons which escapes (e.g.\\ Bianchi et al.\\ 2001; Heckman et al.\\ 2001; Hurwitz, Jelinsky, \\& Dixon 1997; Leitherer et al.\\ 1995).} of UV ionizing photons of $f_{\\mbox{\\tiny esc}} \\simgtr 0.5$.  Given the evidence for a strong outflow in ES1, it would be intriguing to measure the flux of photons below $\\lambda = 912$ \\AA.  However, ES1 is very faint even above \\lya; we estimate that it fades to $I_{\\mbox{\\tiny AB}} > 29$ below the Lyman limit.  We can do only slightly better at high redshifts with more accessible spectra: even in a composite of four of the highest signal--to--noise Keck spectra of galaxies at $z > 4.5$ collected by Spinrad and collaborators (Figure~\\ref{composite}), our measurement of the flux of 900 \\AA\\ photons is consistent with zero at the $1 \\sigma$ level.  This result translates to the coarsely constrained flux ratio $\\fele / \\fnin = 16.7 \\pm 51.9$ ($1 \\sigma$ uncertainty).  To convert this value to the more useful $\\ffif / \\fnin$ ratio, we adopt an empirical correction factor based on the set of fluxes given in SPA01, yielding an effective $\\ffif / \\fnin = 31.4 \\pm 98.2$. Finally, for the same intrinsic Lyman discontinuity of 3 adopted by SPA1, we find an escape fraction of $\\fesc \\simgtr 0.1 \\pm 0.3$.  As we did not correct our initial $\\fele / \\fnin$ ratio for the opacity of the IGM, this value represents a lower limit. Evidently, to satisfactorily constrain the correlation between outflow dynamics and the escape of UV ionizing photons in an individual high redshift galaxy, we will require spectroscopy of a lensed, blue candidate system viewed along the outflow direction."
    },
    "0201/astro-ph0201391.txt": {
        "abstract": "We present spectroscopic and polarimetric observations of a complete, optically unbiased sample of 2Jy radio galaxies at intermediate redshifts ($0.15 < z < 0.7$). These data --- which cover the nuclear regions of the target galaxies --- allow us to quantify for the first time the various components that contribute to the UV excess in the population of powerful, intermediate redshift radio galaxies. We find that, contrary to the results of previous surveys --- which have tended to be biased towards the most luminous and spectacular objects in any redshift range --- the contribution of scattered quasar light to the UV excess is relatively minor in most of the objects in our sample. Only 7 objects (32\\% of the complete sample) show significant polarization in the rest-frame UV, and none of the objects in our sample is polarized in the near-UV at the $P > 10$\\% level. Careful measurement and modelling of our spectra have allowed us to quantify the contributions of other components to the UV excess. We show that nebular continuum (present in all objects at the 3 --- 40\\% level), direct AGN light (significant in 40\\% of objects), and young stellar populations (significant in 15 --- 50\\% of objects) all make  important contributions to the UV continuum in the population of powerful radio galaxies. These results serve to emphasise the multi-component nature of the UV continuum in radio galaxies. The results also point to an interesting link betweeen the optical/UV and far-IR properties of our sample objects, in the sense that the objects with the clearest evidence for optical/UV starburst activity are also the most luminous at far-IR wavelengths. This supports the idea that the cooler dust components in radio galaxies are heated by starbursts rather than by AGN.  %As a by-product of %this work on the continuum properties, we also present %emission line ratios for all the objects in our sample, %which have been measured following continuum modelling and subtraction. ",
        "introduction": "Given the large look back times encompassed by the most distant radio sources, one motivation for studying such objects is their potential use as probes of the formation and evolution of giant early-type galaxies in the early universe. However, all studies aimed at using radio galaxies in this way have to face the problem of distinguishing the effects of the AGN and radio jet activity from genuine signs of galaxy evolution. This problem is particularly acute in the case of studies of the continuum properties. Compared with normal early-type galaxies, powerful radio galaxies can show continuum excesses at both optical/UV (e.g. Lilly \\& Longair 1984, Smith \\& Heckman 1989) and far-IR/sub-mm wavelengths (Golombek et al. 1988, Heckman et al. 1994, Archibald et al. 2001). Therefore, a key issue for these objects is whether these continuum excesses are a consequence of recent star formation which may be linked to evolutionary processes in the early-type host galaxies or, given that these objects contain powerful AGN and radio jets, a direct consequence of the activity.  The presence of a UV excess in the continua of radio galaxies was first demonstrated by the photometric observations of samples of high redshift ($z > 0.5$) radio galaxies in the early 1980's. These observations showed evidence for bluer optical-IR colours than expected for non-evolving or passively evolving elliptical galaxies (e.g. Lilly \\& Longair 1984). %It is important to emphasise, however, that this phenomenon is not confined to the %highest redshift radio galaxies. Optical CCD observations of many lower redshift radio %galaxies also show colours that are bluer than expected for passively-evolving elliptical %galaxies (e.g. Smith \\&  Heckman 1989+earlier work?). Initially, the UV excess was interpreted in terms of bursts of star formation, possibly linked to the evolution of the host galaxies. This interpretation is attractive in the light of morphological studies which show evidence for recent mergers in a large fraction of powerful radio galaxies at low redshifts (Heckman et al. 1986); and merger-induced star formation has been suggested as a possible triggering mechanism (Smith \\& Heckman 1989). %If correct, this interpretation ties in with %heirarchical galaxy evolution models which predict relatively recent evolution in the %general population of early-type via galaxy mergers and associated star formation %activity. Radio galaxies may therefore represent a particular --- post-merger --- phase in %the evolution of giant elliptical galaxies (e.g. ???). Indeed, it has also %been argued that galaxy %mergers may be necessary to fuel the activity and also to create the type of rapidly %spinning black holes which may be required to produce powerful radio jets (Wilson \\& %Colbert 1995). However, given the degree of nuclear and extranuclear activity likely to be present in most powerful radio galaxies, some caution is required in deducing starburst properties purely on the basis of broad-band photometric measurements.  Recognising the potential AGN contribution, an alternative explanation for the UV excess was stimulated by the development of the anisotropy-based unified schemes in the late 1980's (e.g. Barthel 1989). In the frame of such schemes the UV excesses can be explained in terms of light scattered from broad radiation cones of the hidden quasar nuclei (Tadhunter et al. 1988, Fabian 1989). Early polarimetric attempts to test this model proved successful in the sense that they showed the high degrees of linear polarization characteristic of anisotropic scattering in the UV continua of several high redshift radio galaxies (e.g. Tadhunter et al. 1992, Cimatti et al. 1993, Vernet et al. 2001). However, while these observations demonstrate that scattered quasar light is an important component of the UV continuum in {\\it some}  sources, they do not establish the significance of the scattered component in the general population of powerful radio galaxies. Because polarimetric observations of faint objects are difficult, previous studies have tended to be biased towards the brightest, most spectacular objects in a given redshift range. There are also redshift-dependent biases which arise because optical (mostly V-band) observations sample the rest-frame UV in the high redshift objects --- with minimal dilution by the old stellar populations of the host galaxies --- but sample the rest-frame optical in the low redshift objects --- with substantial dilution by the old stellar populations. The importance of this observational selection effect is emphasised by multi-wavelength polarimetric observations of individual sources which show a sharp decline in the measured polarization between the UV and the optical (Tadhunter et al. 1996, Ogle et al. 1997, Tran et al. 1998).  In addition to the scattered component, detailed observations over the last decade have revealed the presence of two further activity-related components which can contribute to the UV excess. These are: the nebular continuum emitted by the extended emission line nebulae (Dickson et al. 1995); and direct AGN light emitted by weak, or partially extinguished, quasars in the nuclei of the galaxies (Shaw et al. 1995). The nebular continuum is likely to be particularly significant in regions where the emission lines have large equivalent widths, including the extended emission line nebulae around powerful radio galaxies. In contrast, the direct AGN component will only be important in the nuclear regions of the sources.  Most recently, events have turned full circle with the spectroscopic detection of young stellar populations in at least some powerful radio galaxies (e.g. Tadhunter et al. 1996, Melnick et al. 1997). The detection of this component is consistent with the early interpretation of the UV excess in terms of starbursts associated with the evolution of the host galaxies (Lilly \\& Longair 1984). Unfortunately, apart from cases in which it dominates the optical continuum (e.g. Miller 1981), the starburst component is notoriously difficult to detect at optical wavelengths. Its presence can be masked by the light of the old stellar populations in the bulges of the host galaxies, by the various activity-related continuum components noted above, and by emission lines which can contaminate the absorption features characteristic of young stars. This is illustrated by the case of 3C321 which shows polarimetric evidence for a significant scattered quasar component, but also shows evidence for a starburst component in the form of a Balmer break and Balmer absorption features (Tadhunter et al. 1996, Robinson et al. 2000). It is notable that the starburst component in 3C321 only came to light through detailed modelling of the optical/UV continuum using a combination of spectrophotometry and spectropolarimetry measurements.  Given the complex circum-nuclear environments of powerful radio galaxies revealed by recent HST imaging studies (e.g. Jackson, Tadhunter \\& Sparks 1998), it is not surprising that no single mechanism is responsible for the UV excess. Observations of individual sources demonstrate the presence of at least four UV-emitting components that can contribute to the UV excess: scattered AGN light, direct AGN light, nebular continuum, and the light of young stellar populations. However, the relative importance of these components, and particularly the importance of any starburst component, is not clear from the previously published data. In this paper we attempt to remedy this situation by combining spectroscopic and polarimetric observations to quantify the contributions of the various UV-emitting components in a complete, optically unbiased sample of powerful 2Jy radio galaxies at intermediate redshifts ($0.15 < z < 0.7$). We also consider the link between the optical/UV signs of star formation activity and the far-IR continuum excess. In a companion paper we report a similar study of a lower redshift sample of 3C radio galaxies ($z < 0.2$: Wills et al. 2002).  Throughout this paper we assume a Hubble constant of $H_0 = 50$ km s$^{-1}$ Mpc$^{-1}$ and a deceleration parameter of $q_0 = 0$. ",
        "conclusions": "The results of our survey confirm the multi-component nature of the UV continuum in powerful radio galaxies. We can quantify the contributions of the various components that contribute to the UV excess in the near-nuclear regions as follows: \\begin{itemize} \\item {\\bf Nebular continuum.} This is present in the spectra of all the objects and contributes 5 -- 40\\% of the UV continuum below the Balmer break. \\item {\\bf Direct AGN light.} Based on the detection of broad permitted lines in the intensity spectra, direct AGN light makes a significant contibution to the UV excess in $\\sim$30 -- 40\\% of the objects in our sample. \\item {\\bf Scattered AGN light.} Polarimetric observations provide evidence that scattered AGN light makes a significant contribution to the UV continuum in 37\\% of the 19 objects in our sample with polarization measurements, but in most cases the scattered component does not appear to dominate the UV excess. \\item {\\bf Starburst component.} Young stellar populations clearly dominate the optical/UV continua in three of the objects in our sample ($\\sim$15\\%) and may make a large contribution to the optical/UV continuum in up to 50\\% of all the sample objects. \\end{itemize}  Of these components, the starburst component clearly warrants further investigation. If, as seems likely, the starbursts were triggered by the same merger events that triggered the radio jet and quasar activity, then detailed studies of the starbursts can provide potentially unique information about the genesis of powerful radio sources, particularly the timescales, the nature of the mergers and the order-of-events.  \\subsection*"
    },
    "0201/hep-ph0201156_arXiv.txt": {
        "abstract": "This report is a brief review of the efforts to explain the nature of non--baryonic dark matter and of the studies devoted to the search for relic particles. Among the different dark matter candidates, special attention is devoted to relic neutralinos, by giving an overview of the recent calculations of its relic abundance and detection rates in a wide variety of supersymmetric schemes.  \\vspace{1pc} ",
        "introduction": "The presence of large amounts of non--luminous components in the Universe has been identified along the years by different means and on different scales: on the galactic scale, the flatness of the rotational curves of many galaxies indicates a dark component which is presumably distributed as a halo around the galaxies \\cite{Rot.curves}; clusters points toward a sizeable contribution of unseen matter distributed around and between galaxies \\cite{Clusters}; more recently, on cosmological scales, the combination of the results on high--redshift supernovae \\cite{Highz.SN} and on the anisotropies of the cosmic microwave background radiation \\cite{CMBR} is pointing toward a flat Universe whose energy density is dominated by a dark vacuum component (cosmological constant, quintessence) together with a sizeable dark component of matter. In terms of the density parameter $\\Omega$, the current view can be summarized as follows \\cite{deBernardis}: the total amount of matter/energy of the Universe is $\\Omega_{\\rm tot} = 1.02 \\pm^{0.06}_{0.05}$, and this is composed of a matter component $\\Omega_{\\rm M} = 0.31 \\pm^{0.13}_{0.12}$ and a vacuum--energy component $\\Omega_{\\Lambda} = 0.71 \\pm {0.11}$. Even though the actual numbers vary a little depending on the priors of the statistical analyses, the clear indication of the latest data is that the Universe is strongly dominated by dark (and unknown) components. In fact the numbers above cannot be reconciled with a Universe made only of standard components: from primordial nucleosynthesis studies, baryons can contribute only at the level of $\\Omega_{\\rm b} = 0.037 \\pm 0.011$ \\cite{Baryons}, while luminous matter is known to provide only a contribution of order $\\Omega_{\\rm lum} \\sim 0.003$ \\cite{Luminous}. We are therefore facing the presence of at least {\\em three dark components} in the Universe: dark baryons, dark non--baryonic matter and dark energy.  The existence of both dark (relativistic or non--relativistic) exotic matter and dark energy asks for extension of the standard model of fundamental interactions, since no known particle or field can explain either of these components.  In this review, we will discuss the current status of the non--baryonic dark matter problem. For reviews on dark baryons and dark energy, see Refs. \\cite{Talk.b,Talk.l}. We will discuss the efforts to explain the amount of dark matter in the Universe, which we summarize as: \\begin{equation} 0.05 \\lsim \\Omega_{\\rm M} h^2 \\lsim 0.3 \\label{eq:DM} \\end{equation} and the studies related to the searches for dark matter particles.  \\vspace{-5pt} ",
        "conclusions": "We have seen that we can identify two main issues in particle dark matter studies: {\\em i)} to explain the observed amount of dark matter in the Universe ($0.05 \\lsim \\Omega_{\\rm M} h^2 \\lsim 0.3$) by finding suitable particle candidates; {\\em ii)} to detect a relic particle. For both of them there appear to be good prospects of success.  As for the candidates, there are many proposed particles which could act as dark matter. Some of these candidates turn out to be quite natural, like {\\em e.g.} the massive neutrino, the axion or the neutralino. Almost all of the proposed candidates can play the role of the dominant dark matter component, although for some of them a non--standard cosmology is required. An important remark is that, from the particle physics point of view, dark matter may naturally be multi--component. A multi--component dark matter scenario offers opportunity for interesting phenomenology not only to the dominant candidate, which would explain the cosmological observation on the $\\Omega_{\\rm M}$ parameter, but also to the sub--dominant candidates, since usually these are the ones which are easier to detect. The detection of a particle which is a relic from the early Universe would be a very important and exciting result.  As for detection, perspectives are good, both for direct and for indirect detection techniques, especially for the most interesting and studied candidate, the neutralino. The possibility to have detectable rates for neutralinos depends on the specific supersymmetric model which is considered, and quite generally it appears simpler to detect a relic neutralino which is a sub--dominant dark matter component.  Nevertheless, there are many supersymmetric schemes where relic neutralinos can provide enough cosmological abundance to explain the observed amount of dark matter, and at the same time they can have detection rates large enough to be accessible to direct, and also to some indirect, detection methods. The positive indication of annual modulation in the detection rate of the DAMA/NaI Collaboration can be interpreted as pointing towards a direct detection of a relic massive particle. This effect is, at the moment, the most compelling indication for a particle dark matter signal. When interpreted as originated from relic neutralinos, the annual modulation effect can be explained in a number of realization of supersymmetry. It is worth noticing that the presence of a signal from dark matter, like the annual modulation effect or signals which could hopefully come in future experiments, can be very important not only for astrophysics and cosmology but also for particle physics, since the need to explain the effect can help in deriving properties of particle physics models and possibly discriminate among different realizations, for instance of supersymmetry.  \\smallskip {\\bf Acknowledgements.} This work was partially supported by the Research Grants of the Italian Ministero dell'Universit\\`a e della Ricerca Scientifica e Tecnologica (MURST) within the {\\sl Astroparticle Physics Project}. I also wish to thank the Organizers for inviting me to deliver this talk at the TAUP 2001 Conference."
    },
    "0201/astro-ph0201149_arXiv.txt": {
        "abstract": "We analyse the non-linear propagation and dissipation of axisymmetric waves in accretion discs using the ZEUS-2D hydrodynamics code.  The waves are numerically resolved in the vertical and radial directions.  Both vertically isothermal and thermally stratified accretion discs are considered.  The waves are generated by means of resonant forcing and several forms of forcing are considered.  Compressional motions are taken to be locally adiabatic ($\\gamma = 5/3$).  Prior to non-linear dissipation, the numerical results are in excellent agreement with the linear theory of wave channelling in predicting the types of modes that are excited, the energy flux by carried by each mode, and the vertical wave energy distribution as a function of radius.  In all cases, waves are excited that propagate on both sides of the resonance (inwards and outwards).  For vertically isothermal discs, non-linear dissipation occurs primarily through shocks that result from the classical steepening of acoustic waves.  For discs that are substantially thermally stratified, wave channelling is the primary mechanism for shock generation.  Wave channelling boosts the Mach number of the wave by vertically confining the wave to a small cool region at the base of the disc atmosphere. In general, outwardly propagating waves with Mach numbers near resonance ${\\cal M}_{\\rm r} \\ga0.01$ undergo shocks within a distance of order the resonance radius. ",
        "introduction": "Gaseous discs play a significant role in the evolution of binary star and planetary systems.  Generally, the tidal forces from stellar or planetary companions distort the discs from an axisymmetric form. However, where resonances occur in the discs, the tidal forces also generate waves that transport energy and angular momentum.  The resulting resonant torques cause orbital evolution of the perturbing objects (e.g. Goldreich \\& Tremaine 1980; Lin \\& Papaloizou 1993; Lubow \\& Artymowicz 1996).  The waves also cause the discs to evolve since they transfer their energy and angular momentum to the disc in the locations where they damp.  Damping may be provided by shocks, radiative damping (Cassen \\& Woolum 1996), or other viscous damping mechanisms.  Waves in discs have also been considered as possible explanations of quasi-periodic variability in the X-ray emission of systems involving accretion discs around black holes (e.g. Kato \\& Fukue 1980; Nowak \\& Wagoner 1991; Kato 2001).  Initial studies of wave propagation approximated the disc as two-dimensional and ignored the effects of vertical structure (perpendicular to the disc plane).  Goldreich \\& Tremaine (1979) developed a two-dimensional linear theory for resonant tidal torques and the associated wave propagation.  Later studies of the non-linear case found that the torques were within a few percent of those predicted by the two-dimensional linear formula (Shu, Yuan \\& Lissauer 1985; Yuan \\& Cassen 1994). For a thin disc that is vertically isothermal and has an isothermal thermodynamic response ($\\gamma = 1$), the only wave excited in the disc has a two-dimensional structure. The wavefronts are perpendicular to the disc plane and the motion is purely horizontal (parallel to the disc plane).  Thus, a two-dimensional treatment is valid in a thin disc if both the vertical structure of the disc and its thermodynamic response are locally isothermal.  However, this is not realistic for most discs, such as those in cataclysmic variables, circumstellar and circumbinary discs around pre-main-sequence stars, and protoplanetary discs.  In such discs, the waves are no longer two-dimensional and the vertical structure of the disc must be taken into account.  A study of three-dimensional wave propagation was made using linearised numerical simulations (Lin, Papaloizou \\& Savonije 1990a,b).  Unfortunately, the limited range of physical parameter space that was covered led to the interpretation that waves in a vertically thermally stratified disc are refracted upwards into the atmosphere where they dissipate via shocks.  Later, it was realised that the linear problem could be solved semi-analytically (Lubow \\& Pringle 1993; Korycansky \\& Pringle 1995; Lubow \\& Ogilvie 1998; Ogilvie \\& Lubow 1999).  These studies found that the propagation of waves in a thin disc is similar to the propagation of electromagnetic waves along a waveguide (Jackson 1962; Landau \\& Lifshitz 1960).  The motions in the disc can be described in terms of modes.  For each mode, the vertical wave structure is determined in detail at each disc radius. The horizontal variations of the mode are treated by means of a WKB approximation in which the horizontal (radial) wavenumber is determined as a function of wave frequency. A flux conservation law determines the mode amplitude as a function of radius.  Depending on the vertical structure of the disc and the thermodynamic response of the gas, the wave energy may be channelled towards the surface of the disc, towards the mid-plane of the disc, or occupy the entire vertical extent as it propagates in radius.  Ideally, one would like to determine the non-linear, non-axisymmetric response of a variety of disc models (vertically isothermal and thermally stratified) to resonant forcing. Unfortunately, this is a formidable numerical problem at present. Instead we begin, in this paper, with an analysis of the non-linear axisymmetric response.  The axisymmetric case offers the major advantage of being expressed as a two-dimensional problem, which is currently accessible numerically. Although axisymmetric waves only transport energy and not angular momentum, they should resemble the response of the disc to low-azimuthal number tidal forcing, as arises in binary star systems. The local physics of the waveguide is determined within a region whose size is a few times the semi-thickness of the disc.  On this scale a non-axisymmetric wave of low azimuthal wavenumber is almost indistinguishable from an axisymmetric wave.  This explains why the local physics of the waveguide can be studied within the shearing-sheet approximation (Lubow \\& Pringle 1993) and why the azimuthal wavenumber appears in the dispersion relation only through a Doppler shift of the wave frequency.  Consequently, in this paper, we study the non-linear propagation and dissipation of axisymmetric, but fully resolved, waves in accretion discs.  The purpose of the paper is threefold.  First, we wish to determine, using non-linear hydrodynamic calculations, whether or not the semi-analytical linear theory provides an accurate description of wave propagation in accretion discs.  Secondly, although linear theory can predict where non-linearity is likely to become important, shocks and the process of energy deposition (and angular momentum deposition with non-axisymmetric waves) cannot be handled.  We aim to determine how and where dissipation occurs through shocks and how accurately this can be predicted from the linear analysis.  Finally, we wish to determine how well non-linear hydrodynamic calculations can model the problem and what resolution is required.  The outline of this paper is as follows.  In Section 2, we briefly review the ZEUS-2D hydrodynamic code that is used to obtain the non-linear results.  In Section 3, we discuss the structure of the equilibrium discs that we model and the grid layout for the ZEUS-2D calculations.  Section 4 briefly reviews the semi-analytical linear method for solving the three-dimensional wave propagation problem, describes our method of wave excitation, and discusses the requirements for convergence of the numerical calculations.  Sections 5 and 6 contain the results for a vertically isothermal disc, and a polytropic disc with a vertically isothermal atmosphere, respectively. Intermediate discs with optical depths not much larger than unity are considered in Section 7.  A summary of the results is contained in Section 8. ",
        "conclusions": "\\subsection{Comparison with linear theory}  We have performed non-linear numerical simulations of the excitation, propagation and dissipation of axisymmetric hydrodynamic waves in an accretion disc.  The disc is Keplerian, vertically resolved and inviscid apart from an artificial numerical bulk viscosity that is required in order to resolve shocks.  The waves are excited by applying a temporally periodic acceleration throughout the computational domain, and are launched resonantly at radii where the forcing frequency coincides with a natural oscillation frequency of the disc.  We have examined Lindblad resonances, vertical resonances and hybrid vertical/Lindblad resonances.  In each case the waves propagate radially through the disc away from the resonance, changing continuously in vertical structure as they do so, and ultimately dissipate.  This mechanism of generating the waves is directly comparable with the way that non-axisymmetric waves are excited in discs by tidal forcing from an orbiting companion in a binary or protoplanetary system.  A low-mass companion such as a sub-Jovian planet orbiting a star exerts its tidal influence on the disc mainly through launching \\fe-mode waves at the Lindblad resonances (Goldreich \\& Tremaine 1980). Lindblad resonances are also important in systems of larger mass ratio, such as circular orbit binary star systems of extreme mass ratio (Lin \\& Papaloizou 1979) and eccentric orbit binaries (Artymowicz \\& Lubow 1994).  Related resonances play a role in the growth of the eccentric disc mode in superhump binaries  (Lubow 1991) and possibly binaries with brown dwarf secondaries (Papaloizou, Nelson \\& Masset 2001). Vertical resonances are an important aspect of the tidal interaction in close binary stars (Lubow 1981; Ogilvie 2002), where the Lindblad resonances are all excluded from the disc.  The corrugation-mode resonances we have studied are the axisymmetric equivalent of the bending-mode resonances observed in Saturn's rings (Shu, Cuzzi \\& Lissauer 1983).  Wherever the waves are linear, our results are in excellent and detailed agreement with the linear theory developed by Lubow \\& Pringle (1993), Korycansky \\& Pringle (1995), Ogilvie (1998), Lubow \\& Ogilvie (1998) and Ogilvie \\& Lubow (1999).  Elements of this theory can also be found in the work of Loska (1986) and Okazaki, Kato \\& Fukue (1987).  The disc acts as a waveguide that allows a variety of wave modes to propagate in the radial direction.  The modes can be classified similarly to stellar oscillations, and we have discussed all four classes (f, p, r and g) to some extent in this paper.  Each mode has a dispersion relation that connects the wave frequency and the radial wavenumber at every radius.  In general, waves of a given frequency can propagate only in restricted radial intervals, bounded by turning points where the wavenumber vanishes.  The turning points correspond to the resonant radii where waves are launched by an applied periodic forcing.  The linear theory is confirmed by the simulations in a number of respects.  First, the concept of the disc as a waveguide that supports discrete propagating modes is validated.  The waves are vertically confined by the continuous stratification of the disc, not by zero-density surfaces, artificial boundaries or other discontinuities. The solutions exhibit the same `separation of variables' that also features in the linear analysis.  That is, the waveform is an oscillatory function of radius (and time) multiplied by a vertical profile that changes slowly with radius.  The fact that the velocity amplitudes of the linear solutions may formally diverge high in the atmosphere does not invalidate the rest of the solution when the energy density of the wave is vertically confined in the bulk of the disc.  Secondly, the energy fluxes imparted to the different wave modes through launching at different types of resonance are in excellent agreement with the predictions of linear theory presented in Appendix~B.  This analysis of resonant wave launching is closely related to the well-known resonant torque formula of Goldreich \\& Tremaine (1979) and to the analysis of vertical resonances by Lubow (1981).  Finally, the detailed waveforms predicted by linear theory are confirmed in the simulations (see, e.g., Figure \\ref{refraction} and the next subsection), even when the disc is not especially thin.  \\begin{figure*} \\vspace{18.0truecm} \\caption{\\label{refraction} Comparison of the simulations with linear theory, and comparison of wave channelling with acoustic refraction.  In the upper grey-scale figure, we plot a snapshot of the radial velocity multiplied by the square root of the density, $v_{\\rm r}\\sqrt{\\rho}$, from a simulation involving an \\fe mode in a polytropic disc that joins smoothly on to a low-mass isothermal atmosphere ($H/r=0.1$, $\\tau=25000$). (The square of this quantity is the radial part of the instantaneous wave energy density, c.f.\\ equation 5). The amplitude of the wave is ${\\cal M}_{\\rm r}=0.001$ and the grid resolution is $N_r\\times N_{\\theta}=1000\\times 183$.  In the lower grey-scale figure, we plot the corresponding semi-analytical prediction of the linear theory of wave channelling, ignoring the r~mode and using the WKB approximation in the radial direction (which results in a mild singularity at the Lindblad resonance).  The agreement is excellent until dissipation occurs.  Over the upper grey-scale, we trace the rays that would result from the refraction of an initially vertical wavefront at the resonant radius.  The refracted wavefronts, which are orthogonal to the rays, would rapidly become severely tilted and the wave would be predicted to propagate almost vertically into the isothermal atmosphere.  However, the simulation shows that the wavefronts in fact remain nearly vertical as the wave propagates horizontally along the base of the disc atmosphere (shown by the dashed line, $z=H$).  The \\fe mode is generated throughout the thickness of the disc at $r=1$ but its energy rises from within the polytropic layer owing to wave channelling and becomes confined near the base of the atmosphere.  The wave propagates {\\it without loss of flux\\/} to $r\\approx 2.2$ (Figure \\ref{PF_ep0.10}, dot-dashed line, left panel) before most of its energy is dissipated in two regions at the base of the atmosphere (white contour lines at $r\\approx 2.2$). The dissipation in this example is partly numerical, as suggested by Figure \\ref{PF_ep0.10}.  If we were able to increase the resolution further, the wave would propagate a greater distance (c.f.\\ different resolutions in the left panel of Figure \\ref{PF_ep0.10}) and might agree even more closely with the linear prediction. } \\end{figure*}  \\subsection{Wave channelling versus refraction}  The continuous change of the vertical profile of a wave as it propagates radially is a process we have termed `wave channelling' (Lubow \\& Ogilvie 1998).  The energy density of the wave can be channelled either towards the surfaces of the disc or towards the mid-plane as it propagates away from its site of launching.  The results of this paper provide ample evidence for both types of wave channelling.  In particular, the f and p~modes in a disc with a vertical temperature gradient are channelled towards the surfaces (e.g. Figures \\ref{PF_ep0.10_cont}, \\ref{PP_ep0.10_cont} and \\ref{PT_ep0.10_cont}), while the r~modes in a disc with a vertical entropy gradient are channelled towards the mid-plane (e.g. Figures \\ref{IF_ep0.10_cont} and \\ref{IT_ep0.10_cont}) as predicted by Lubow \\& Pringle (1993).  The wave energy of the \\fe mode, which is the principal mode launched at a Lindblad resonance, rises towards the surfaces of a thermally stratified disc.  Close to the resonance, the \\fe mode occupies the full vertical extent of the disc and behaves compressibly.  As the mode propagates away from the resonance, its energy within the disc midplane region rises and becomes confined (or channelled) to a layer at the base of the disc atmosphere. The extent of channelling depends on the degree of thermal stratification. In this regime, the wave becomes a surface gravity wave and behaves incompressibly (Ogilvie 1998; Lubow \\& Ogilvie 1998). {\\it The process by which this occurs is wave channelling and not acoustic refraction.} This is demonstrated as follows.  If refraction were responsible (e.g.\\ Lin et~al. 1990a,b), the wave would be directed upwards into the isothermal atmosphere after travelling a radial distance of order $H$.  This is not observed in the simulations.  Plots of the wave energy for low-amplitude waves in polytropic discs (Figures \\ref{PF_ep0.10_cont}, \\ref{PP_ep0.10_cont} and \\ref{PT_ep0.10_cont}) clearly show that the waves propagate along the base of the atmosphere before dissipating in shocks.  Propagation does not switch from horizontal to vertical, as would be expected if simple refraction were involved.  In fact, the \\fe mode is launched even in an incompressible fluid with a vertical energy distribution that is similar to the compressible case.  Consequently, the process of energy concentration cannot be due to acoustic refraction.  Figure \\ref{refraction} shows an explicit comparison between the results of the simulations and the predictions of acoustic refraction. In order to show both the wavefronts and the vertical extent of the wave, we plot a snapshot of the radial velocity, normalized by the RMS velocity, and multiplied by the mean wave energy density, from a simulation involving an \\fe mode in a polytropic disc.  On top of the grey-scale, we trace the rays that would result from the refraction of an initially vertical wavefront at the resonant radius.  The refracted wavefronts, which are orthogonal to the rays, would rapidly become severely tilted and the wave would be predicted to propagate almost vertically into the isothermal atmosphere.  However, the simulations show that the wavefronts in fact remain nearly vertical as the wave propagates horizontally along the base of the disc atmosphere.  The wave propagates {\\it without loss of flux\\/} to $r\\approx 2.2$ before most of its energy is dissipated in two regions at the base of the atmosphere (see Figure \\ref{PF_ep0.10} for the radial energy flux). Figure \\ref{refraction} demonstrates that the linear theory of wave channelling accurately predicts the propagation of the wave until non-linear dissipation occurs.  In summary, our non-linear hydrodynamic calculations clearly show that for axisymmetric waves the disc acts as a waveguide and that the behaviour of the waves is determined primarily by wave channelling and not refraction.  Explanations based on simple refraction are incorrect.  Ray tracing, if applied correctly, does offer a valid description of modes with a short vertical wavelength, i.e. those of large vertical mode number (see Figure 11 of Lubow \\& Pringle 1993).  However, such modes are unlikely to be excited by tidal resonant forcing and the \\fe mode, in particular, is vertically evanescent and has a vertical mode number of zero. As mentioned in Section 1, non-axisymmetric waves of low azimuthal wavenumber $m\\lsim r/H$ are almost indistinguishable from axisymmetric waves on the scale of a few $H$. We therefore expect that such waves will also propagate as if through a waveguide and be subject to similar wave channelling in agreement with linear theory.  Several papers (e.g.\\ Lin et al.\\ 1990a,b; Terquem 2001) have interpreted the behaviour of low-$m$ (e.g.\\ $m=2$) non-axisymmetric waves as being due to refraction.  Given our results for $m=0$ waves, we conclude that the dominant behaviour of low-$m$ non-axisymmetric waves is almost certainly determined by wave channelling and not refraction.   \\subsection{Dissipation of waves}  In situations where waves in discs are excited by tidal forcing from an orbiting companion, the site at which a wave ultimately dissipates is of some importance, because the energy and angular momentum carried by the wave are transferred to the disc there.  The simulations provide valuable information on the location and means of wave dissipation, matters about which we were previously able only to speculate (e.g. Lubow \\& Ogilvie 1998).  In a nearly inviscid disc, dissipation can occur only if the wave develops very large velocity gradients.  One way to achieve this is the classical non-linear steepening of an acoustic wave in a gas with $\\gamma>1$ (e.g. Lighthill 1978).  The crest of the wave travels faster than the trough and the wave steepens as it propagates until shocks form.  This is the principal effect leading to the dissipation of the \\fe mode in a vertically isothermal disc with $\\gamma=5/3$ (Figure \\ref{IF_ep0.10_peak}).  The situation is not entirely straightforward because the wave Mach number is not uniform across the wavefront.  In addition, the steepening may be accelerated by the gradual increase of wave amplitude as the wave propagates outwards into material of lower surface density while conserving its energy flux.  On the other hand, the steepening may be temporarily postponed by the dispersive character of the wave close to the resonance.  The r~modes in the vertically isothermal disc are channelled into an increasingly thin layer near the mid-plane and simultaneously develop very short radial wavelengths.  These modes are therefore susceptible to viscous damping even in a nearly inviscid disc.  In our simulations, these modes develop scales shorter than the grid spacing and are ultimately lost.  In the polytropic disc, classical wave steepening does not occur in any simple sense because none of the waves behaves like a plane acoustic wave.  The \\fe mode, in particular, is highly dispersive, which tends to resist non-linear steepening.  It undergoes rapid wave channelling to the base of the isothermal atmosphere and behaves like a surface gravity wave.  The wave channelling enhances the amplitude of the wave by concentrating its energy into a smaller region of space.  In earlier work (Lubow \\& Ogilvie 1998; Ogilvie \\& Lubow 1999) we showed that this effect is typically sufficient to amplify the wave to sonic velocities where steepening and dissipation are unavoidable. At the same time, the radial and vertical length-scales of the wave also decrease as it propagates outwards (Figure \\ref{refraction}), making linear viscous damping a possible source of dissipation. Because wave channelling occurs over a distance that is almost independent of the thickness of the disc, the dissipation of waves with the same initial Mach number near the reasonance occurs at the same radius regardless of the disc's thickness.  We have performed calculations with discs whose thicknesses vary by a factor of 4 to demonstrate this effect (Section \\ref{DependenceOnThickness} and Table \\ref{table1}). The results of the simulations suggest that the \\fe mode in a polytropic disc dissipates energy at the base of the isothermal layer at a Mach number of approximately $0.2$ (Figure \\ref{PF_ep0.10_peak_a}), which it acquires through the effect of wave channelling.  The wave energy does not propagate vertically within the isothermal atmosphere, contrary to the suggestion of Papaloizou \\& Lin (1995).  \\subsection{Outlook}  In future we hope to address some questions that remain unanswered in this paper.  In particular, a more detailed study of the competition between linear dispersion and non-linear steepening of waves in discs would be valuable.  Some important wave phenomena will require non-axisymmetric simulations.  For example, a Keplerian disc supports slowly varying $m=1$ density and bending waves that effect an eccentric distortion and a warping of the disc, respectively.  Unlike those studied in this paper, these waves can propagate over long distances without experiencing wave channelling.  Finally, attention should be given to studying wave propagation in possibly more realistic models such as fully turbulent discs and layered discs (Gammie 1996).  These challenges provide ample opportunities for further developments."
    },
    "0201/astro-ph0201180_arXiv.txt": {
        "abstract": "We present the complete submillimeter data for the Canada-UK Deep Submillimeter Survey (CUDSS) 3\\h field.  The observations were taken with the Submillimeter Common-User Bolometer Array (SCUBA) on the James Clerk Maxwell Telescope (JCMT) on Mauna Kea.     The 3\\h field is  one of two main fields in our  survey and covers 60 square arc-minutes to a 3$\\sigma$ depth of $\\sim$ 3 mJy.  In this field we have detected 27 sources above 3$\\sigma$, and 15 above 3.5$\\sigma$.  We assume the source counts follow the form $N(S) {\\propto} S^{-\\alpha}$ and measure $\\alpha$ = 3.3$^{+1.4}_{-1.0}$. This is in good agreement with previous studies and further supports our claim  \\citep{eal00} that SCUBA sources brighter than 3 mJy produce $\\sim$20\\% of the 850$\\mu$m background energy.   Using preliminary ISO 15 $\\mu$m maps and VLA 1.4 GHz data we have identified counterparts for six objects and have marginal detections at 450$\\mu$m for two additional sources.     With this information we estimate a median redshift for the sample of 2.0$\\pm$0.5, with $\\sim$10\\% lying at $z<$ 1. We have measured the angular clustering of  \\s8$>$ 3 mJy sources using the source catalogues from the CUDSS two main fields, the 3\\h and 14\\h fields, and find a marginal detection of clustering, primarily from the 14\\h field, of $\\omega(\\theta)=4.4\\pm2.9  \\theta^{-0.8}$.  This is consistent with clustering at least as strong as that seen for the Lyman-break galaxy population and the Extremely Red Objects.  Since SCUBA sources are selected over a broader range in redshifts than these two populations the strength of the true spatial clustering is expected to be correspondingly stronger. ",
        "introduction": "Over the last decade there have been great steps forward in our understanding of the formation and early evolution of galaxies.  There are currently  two  general, distinct theories of massive galaxy formation, though the true picture is likely a combination of the two.   In the first, galaxies form  over a range of redshift, from the gradual hierarchical merging of smaller aggregates \\citep{bau96,kau98}. In this picture galaxy formation is an ongoing  process characterised  by star formation rates of moderate magnitude. In the second scenario galaxies form at high-redshift on short timescales from the collapse of a single  object and undergo one massive burst of star formation \\citep{egg62}.  They then evolve passively to galaxies of the  present-day.  The observational picture is still  somewhat confused. Optical studies have found that the  luminosity density of the universe increases out to z$\\sim$1 \\citep{lil96,mad98,hog98} and, according to the observations of the   Lyman-break galaxy population (LBG), does not decrease to at least   z$\\sim$4 \\citep{ste99}. However,  the star formation rates in individual LBGs of a few tens of solar masses per year \\citep{ste96}, though large compared to local starbursts, are too moderate  to form an elliptical galaxy on  a dynamical timescale of $\\sim$10$^8$ years, and suggest gradual, hierarchical  formation. The hierarchical model  is further supported by the  increase in the  rate of galaxy-galaxy interactions with redshift \\citep{pat01}.   On the other hand, the spheroids, which contain $\\sim$2/3 of the stars in the universe \\citep{fuk98}  still appear to be old at z$\\sim$1 \\citep{zep97,cim99,sco00,mor00}.  The homogeneous nature of their  stellar populations today \\citep{bow92} imply formation over a  short timescale and at high redshift. However, until recently, no high-redshift object with star formation rates large enough to form a massive spheroid in a dynamical timescale had been seen.  The deep submillimeter surveys of the last five years, \\citep{sma97,hug98,bar98,eal99} with the Submillimeter Common-User Bolometer Array (SCUBA) on the James Clerk Maxwell Telescope (JCMT), have uncovered just such a population.  \\  The results of these deep SCUBA surveys have been exciting and in general agreement with each other.  The population revealed by the SCUBA surveys covers a broad range of redshift, with a median  redshift of 2 $<$ z $<$ 3 \\citep{eal00,dunl01}.   Many of the  submillimeter detections which have secure optical/near-infrared  (NIR) counterparts show disturbed morphology or multiple-components suggestive of galaxy mergers \\citep{lil99,ivi00}. These objects have  spectral energy distributions broadly similar  to today's ultra-luminous infrared galaxy (ULIRG) population.   In the  local universe the ULIRGs are the most luminous galaxies and emit the bulk of their energy at FIR wavelengths.  The FIR emission is from dust which is currently thought to be heated by young stars  \\citep{lut99}.The ULIRGs are primarily the result of mergers and result in objects with surface-brightness profiles of elliptical galaxies (see \\citet{san96} for a review).  Though the dust temperature of the SCUBA sources  is very poorly known, if we assume a temperature similar to the local ULIRGS then  these objects are extremely luminous, with  bolometric luminosities of  10$^{12-13}$L$_{\\odot}$. Radio and CO observations of two SCUBA sources \\citep{ivi01} have detected possible extended emission which is in marked difference to  the compact nature of local ULIRGs. It is still unclear whether the majority of these objects are powered by star formation or active galactic nuclei (AGN). Evidence is  mounting through X-ray and optical emission line measurements that, although AGN are present in a small fraction of sources, star formation is, by far, the dominant process \\citep{ivi00,fab00,bar01}.  Given this, the ULIRGs must be forming stars at unprecedented rates of hundreds to  thousands of solar masses per year \\citep{ivi00,eal99}.   These high star formation rates, together with the contribution that these objects make to the total extragalactic background, showing this is a cosmologically-significant population \\citep{eal00}, makes it hard to avoid the conclusion that these objects are elliptical galaxies being seen during their initial burst of star formation. \\     Analysis of the spatial clustering of different populations can provide clues to their evolutionary connections. Recent measurements  of the clustering of two other populations of high-redshift galaxies,  the LBGs and Extremely Red Objects (EROs), \\citep{gia98,gia01,dad00,dad01} have yielded surprising results.  Studies of the z $<$ 1 universe \\citep{lef96,carl00} have found  the clustering strength of galaxies to decrease with increasing redshift, as  expected in a scenario where structure forms through gravitational instabilities.  However, the LBGs ($z \\sim$3) and the EROs ($z \\gtsim $1) are very strongly  clustered. With hindsight this result is in agreement with the prediction of \\citet{kai84}  that the highest peaks in the density field of the early universe  should be strongly clustered.  At the redshifts of LBGs and EROs the universe was much younger and there had been less time for gravitational collapse. These objects are therefore probably the result of the collapse of the rare high peaks in the density field.  As SCUBA sources are even rarer than the LBGs,  and  based on their star formation rates perhaps more massive, they would also be expected to show clustering. \\   \\citet{pea00} investigated the underlying structure in a submillimeter map of the Hubble Deep Field \\citep{hug98}, after removing all discrete  sources above 2 mJy, and found no significant clustering of the underlying flux.  However, the HDF area is small and much larger areas are needed to investigate the clustering of submillimeter sources. Of the current deep SCUBA surveys there are  two blank-field surveys of significant size and with which a clustering measurement may be made: the ``8 mJy survey'' \\citep{sco01},  and our own, the Canada-UK Deep Submillimeter Survey  (CUDSS).  The ``8 mJy survey''  has detected a clustering signal for $\\s8 > $ 8 mJy sources over an area of 260 arcmin$^2$. Our survey  covers 100 arcmin$^2$ and reaches a deeper depth of 3 mJy. \\   This paper is the sixth of a series of papers on the CUDSS project and contains the complete submillimeter data of our 3\\h field. The submillimeter survey is now complete and the final catalogue contains 50 sources, 27 of which have been detected in the 3\\h field. Paper I \\citep{eal99} introduces the survey and  initial detections; paper II \\citep{lil99} discusses the first optical identifications;  paper III \\citep{gea00} discusses the multi-wavelength properties of a particularly interesting and bright source, 14-A;  paper IV \\citep{eal00} presents the nearly complete 14$^h$ field submillimeter sample and  discusses the mid-IR and radio properties of the sources;  paper V \\citep{web01} investigates the relationship between SCUBA sources and LBG galaxies in the CUDSS fields; and papers VII and VIII (Clements et al., in preparation; Webb et al., in preparation) will discuss the optical and near-IR properties of the entire sample.    \\  This paper is laid out as follows:  \\S 2 describes the submillimeter observations, \\S 3  discusses the data reduction and analysis techniques, \\S 4 presents the source catalogue,  \\S 5  discusses the radio and ISO data, in \\S 6 we discuss individual sources, in \\S 7 we present the source counts, in \\S 8 the clustering analysis is performed and the implications of these results are discussed in \\S 9. ",
        "conclusions": "We have used SCUBA on the JCMT to map 60 square arc-minutes of the CFRS 3\\h field.   We have detected 27 sources,  bringing the final number of objects at \\s8$\\gtsim$ 3 mJy detected in the CUDSS to 50.  We have found the following results:  \\begin{enumerate} \\item{For the differential source counts ($N(S){\\propto}S^{-\\alpha}$) we measure $\\alpha$ = 3.3$^{+1.4}_{-1.0}$ which is in excellent agreement with other studies.  Down to 3 mJy these objects are responsible for $\\sim$20\\% of the 850$\\mu$m background energy} \\item{We have used preliminary ISO 15$\\mu$m data, VLA 1.4 GHz observations, and SCUBA 450$\\mu$m maps to identify counterparts of the 850$\\mu$m sources.  Using spectroscopy from the CFRS and the radio-to-submillimeter redshift estimator \\citep{car99,dun00} we have estimated the mean redshift to be 2.0$\\pm$0.5 with  10\\% of the objects below $z <$ 1.} \\item{We have measured the angular clustering of \\s8$>$ 3 mJy sources using the complete CUDSS 3\\h and 14\\h catalogues.  We find $\\omega(\\theta)=4.4\\pm2.9 \\theta^{-0.8}$.  This is as strong as the angular clustering measured for LBGs and EROs, and the spatial clustering will be even stronger due to the broad redshift range of SCUBA sources compared to  LBGs and EROs} \\end{enumerate}  {\\it Acknowledgments} We are grateful to the many members of the staff of the Joint Astronomy Centre who have helped us with this project.  Research by Simon Lilly is supported by the National Sciences and Engineering Council of Canada and by the Canadian Institute of Advanced Research. Research by Tracy Webb is supported by the National Sciences and Engineering Council of Canada and by the Canadian National Research Council.  Research by Stephen Eales, David Clements, Loretta Dunne and Walter Gear is supported by the Particle Physics and Astronomy Research Council.  The JCMT is operated by the Joint Astronomy Centre on behalf of the UK Particle Physics and Astronomy Research Council, the Netherlands Organization for Scientific Research and the Canadian National Research Council.   We also thank Ray Carlberg for many helpful discussions.   \\"
    },
    "0201/astro-ph0201463_arXiv.txt": {
        "abstract": "Supernova--driven outflows from early galaxies may have had a large impact on the kinetic and chemical properties of the intergalactic medium (IGM).  We use three--dimensional Monte Carlo cosmological realizations of a simple linear peaks model to track the time evolution of such metal--enriched outflows and their feedback on galaxy formation.  We find that at most 30\\% of the IGM by volume is enriched to values above $10^{-3} Z_\\odot$ in models that only include objects that cool by atomic transitions. The majority of enrichment occurs relatively early ($5 \\lsim z \\lsim 12$) and leads to a mass-averaged cosmological metallicity between $10^{-3} Z_\\odot$ and $10^{-1.5} Z_\\odot$. The inclusion of Population III objects that cool through ${\\rm H}_2$ line emission has only a minor impact on these results: increasing the mean metallicity and filling factor by at most a factor of 1.4, and moving the dawn of the enrichment epoch to $z \\approx 14$ at the earliest. Thus enrichment by outflowing galaxies is likely to have been incomplete and inhomogeneous, biased to the areas near the starbursting galaxies themselves. Models with a $10\\%$ star formation efficiency can satisfactorily reproduce the nearly constant ($2\\leq z\\leq 5$, $Z\\approx 3.5\\times 10^{-4}\\,Z_\\odot$) metallicity of the low column density Ly$\\alpha$ forest derived by Songaila (2001), an effect of the decreasing efficiency of metal loss from larger galaxies. Finally, we show that IGM enrichment is intimately tied to the ram-pressure stripping of baryons from neighboring perturbations. This results in the suppression of at least 20\\% of the dwarf galaxies in the mass range $\\sim 3\\times 10^8-3\\times 10^9\\,M_\\odot$ in all models with filling factors greater than 2\\%, and an overall suppression of $\\sim 50\\%$ of dwarf galaxies in the most observationally-favored model. ",
        "introduction": "Recent QSO absorption line observations have shown that the intergalactic medium (IGM) is polluted with heavy elements at intermediate redshifts (Songaila \\& Cowie 1996). From such measurements of column density ratios $N_{\\nCIV}/N_{\\nHI}$, Hellsten et al.\\ (1997) and Rauch, Haehnelt, \\& Steinmetz (1997) concluded that typically [C/H]~$ \\simeq -2.5$ at $z \\simeq 3$, with an order of magnitude dispersion about this mean value.\\footnote{In the usual notation, [C/H] = log (C/H) $-$ log (C/H)$_{\\odot}$.}  These values, however, refer to overdense regions of the universe, traced by \\Lya\\ clouds with column densities in excess of $\\log N_\\nHI = 14.5$.  The presence of metals has more recently been assessed in clouds in which $\\log N_\\nHI < 14.0$ as reviewed by Pettini \\etal (2001). At these low optical depths, statistical techniques to extend the search for highly ionized species such as \\CIV and \\OVI must be applied. The results show that {\\it i)} unrecognized weak \\CIV systems must be present in order to reproduce the full \\CIV optical depth (Ellison \\etal 2000), and {\\it ii)} that metals, as traced by \\OVI are present in a gas with a density lower than that of the mean IGM (Schaye \\etal 2000). Very recently, Songaila (2001) has been able to trace the IGM metallicity evolution in systems with $\\log \\CIV>12$ and conclude that a minimum metallicity $Z\\approx 3.5\\times 10^{-4} Z_\\odot$ is already in place at $z=5$.  Although these techniques help to extend QSO absorption studies to underdense regions of the IGM, present observations are only able to place a lower limit on the total volume filling factor of metals. Current measurements, combined with numerical simulations, indicate that metals associated with $\\log N_\\nHI\\simlt 14.2$ filaments fill $\\simgt 3\\%$ of intergalactic space, including areas far away from the high overdensity peaks where galaxies form (Madau, Ferrara, \\& Rees 2001, hereafter MFR).  This suggests that metal pollution occurred relatively early, resulting in a more uniform distribution and enriching vast regions of intergalactic space.  This allows the \\Lya\\ forest to be hydrodynamically `cold' at low redshifts, as intergalactic baryons have time to relax again under the influence of dark-matter gravity.  Note that the presence of high-redshift metals is of great observational importantance, as the measurement of metal lines in $z \\gsim 6$ quasars may also serve as a probe of reionzation (Oh 2002).   These observations prompted MFR to suggest high-redshift ($z \\approx 10$) galaxy outflows as a mechanism for IGM enrichment.  This study could not determine the metal filling factor produced in such a scenario, however, as it was focused on the evolution of typical objects at a single mass-scale and formation redshift.  Similar outflow models have been proposed by Nath \\& Chiba (1995) and Scannapieco \\& Broadhurst (2001, hereafter SB), but primarily motivated by the chemical and thermal properties of the X-ray emitting gas in galaxy clusters.  While the latter of these studies included a range of galaxy masses and was able to make some estimates as to the total filling factor of metals, these results were fairly crude as the study was focused on the properties of individual galaxies.  Aguirre et al.\\ (2001a) and Aguirre et al.\\ (2001b) studied IGM metal enrichment by superimposing an outflow model on numerical simulations that did not include SN-driven winds, but were only able to constrain the contribution from late-forming ($z \\lsim 6$) and relatively large ($M \\gsim 10^{8.5} M_\\odot$) objects. Cen \\& Ostriker (1999) studied metal enrichment in even lower resolution smoothed particle hydrodynamic (SPH) simulations with a dark matter particle mass of $8.6 \\times 10^8 M_\\odot.$ Gnedin \\& Ostriker (1997) studied the relationship between reionization and early metal enrichment in high-resolution simulations, but did not adequately follow supernova explosions.  Finally, Thacker, Scannapieco, \\& Davis (2002) were able to estimate the filling factor of outflows at $z \\geq 4$ purely in the context of high-resolution SPH simulations with a dark matter particle mass of $2.5 \\times 10^6 M_\\odot$, but were not able to examine its dependence on model parameters due to the high computational cost of this approach.  Early-enrichment scenarios also have important implications for the thermal and velocity structure of the IGM, as first studied in Tegmark, Silk, \\& Evrard (1993) and Voit (1996) (see also Cen and Bryan 2000). The resulting feedback on galaxy formation was first examined in Scannapieco, Ferrara, \\& Broadhurst (2000), SB, and Scannapieco, Thacker, \\& Davis (2001). The nature of this effect is twofold: an impinging wind may shock-heat the gas of a nearby perturbation to above the virial temperature, thereby mechanically evaporating the gas, or the gas may be accelerated to above the escape velocity and stripped from the perturbation entirely.  The latter channel is considerably more effective, because shock-heated clouds that are too large to be stripped are able to radiatively cool within a sound crossing time, thus limiting evaporation.  Note that this type of feedback is fundamentally different from the one commonly adopted in galaxy formation models, in which hot gas is produced by supernovae in the parent galaxy.  In this paper we return to the issues of enrichment and feedback, adopting a more complete approach that combines the detailed modeling of a typical object as in MFR, with the more general spatially dependent modeling described in SB.  In this way we are able place constraints on the overall metal filling factor produced as well as investigate the link between cosmic metal enrichment and the feedback from outflows on galaxy formation.  The structure of the paper is as follows.  In \\S2 and \\S3 we describe our semi-analytical simulations of galaxy formation with feedback and IGM enrichment.  In \\S4 we summarize the results of these simulations and the constraints they place of the fraction of the universe impacted by outflows; conclusions are given in \\S5. ",
        "conclusions": "In this work we have studied the metal enrichment of the IGM by outflows in a $\\Lambda$CDM model of structure formation and its feedback on the formation of galaxies.  Adopting a linear peaks model of the spatial distribution of forming objects, and a detailed one-dimensional model of wind propagation, we have determined the overall filling factor as a function of redshift and its relationship with the baryonic stripping of protogalaxies.  While the star formation efficiency of high-redshift galaxies is largely unknown, we are nevertheless able to place useful constraints on the filling factor, enrichment redshift, and overall mass-averaged metallicity in such models.  Choosing a range of star formation efficiencies between $f_\\star = 0.50$ and $f_\\star = 0.01$, we find that at least $3\\%$ and at most $30\\%$ of the IGM is enriched to a level exceeding $10^{-3} Z_\\odot$ by redshift $z=3$.  In all cases, the majority of this enrichment occurs relatively early, $5 \\lsim z \\lsim 12 $, and leads to mass-averaged cosmic metallicities that range from $0.001 Z_\\odot$ to $0.05 Z_\\odot$, for star formation efficiencies $0.01 < f_\\star < 0.5$, respectively.  The mass-averaged metallicity scales roughly linearly with this quantity: $Z \\approx 0.1 f_\\star Z_\\odot.$  Our model can satisfactorily reproduce the constant ($Z\\approx 3.5 \\times 10^{-4} Z_\\odot$) metal enrichment of the low column density Ly$\\alpha$ forest up to $z=5$ derived by Songaila (2001), which is likely to be caused by the decreasing efficiency of metal loss from larger galaxies. This comparison strongly favors star formation efficiencies in a narrow range around 10\\%, essentially excluding the $f_\\star = 0.5$ and $f_\\star = 0.01$ models.  As the formation of stars in Pop III objects is relatively inefficient, the inclusion of these objects has only a secondary effect on our results: increasing the mass-averaged metallicity and filling factors by at most a factor of 1.4, and moving the dawn of the enrichment epoch to $z \\approx 14$ at the earliest.  While all the models studied display suppression of galaxy formation due to outflows ram-pressure stripping the gas out of pre-virialized protogalaxies, this mechanism has only a minor impact on the overall filling factor as it occurs only in the densest and most polluted regions of space.  Nevertheless, after fixing $f_m$, a general relationship between the filling factor and the suppression factor of galaxies exists, at all $f_\\star$ values.  All models and redshifts at which $2\\%$ of the IGM is enriched show a greater than $20\\%$ suppression of galaxies.  In the case that is most consistent with QSO observations, in fact, half the galaxies are suppressed due to baryonic stripping.  Thus the relative quiescence of the Ly$\\alpha$ forest at lower redshifts is likely to belie a violent epoch of early outflows and enrichment."
    },
    "0201/astro-ph0201239_arXiv.txt": {
        "abstract": "The evolution of helium stars with masses of 1.5 -- 6.7~\\msun\\ in binary systems with a 1.4~\\msun\\ neutron-star companion is presented. Such systems are assumed to be the remnants of Be/X-ray binaries with B-star masses in the range of 8 -- 20~\\msun\\ which underwent a case B or case C mass transfer and survived the common-envelope and spiral-in process. The orbital period is chosen such that the helium star fills its Roche lobe before the ignition of carbon in the centre. We distinguish case BA (in which mass transfer is initiated during helium core burning) from case BB (onset of Roche-lobe overflow occurs after helium core burning is terminated, but before the ignition of carbon). We found that the remnants of case BA mass transfer from 1.5 -- 2.9~\\msun\\ helium stars are heavy CO white dwarfs. This implies that a star initially as massive as 12~\\msun\\ is able to become a white dwarf. CO white dwarfs are also produced from case BB mass transfer from 1.5 -- 1.8~\\msun\\ helium stars, while ONe white dwarfs are formed from 2.1 -- 2.5~\\msun helium stars. Case BB mass transfer from more-massive helium stars with a neutron-star companion will produce a double neutron-star binary. We are able to distinguish the progenitors of type Ib supernovae (as the high-mass helium stars or systems in wide orbits) from those of type Ic supernovae (as the lower-mass helium stars or systems in close orbits). Finally, we derive a \"zone of avoidance\" in the helium star mass vs. initial orbital period diagram for producing neutron stars from helium stars. ",
        "introduction": "\\label{helium:sec:intro}  A binary pulsar with a neutron-star or a heavy white-dwarf companion has long been considered to originate from a helium star in a binary system with a neutron-star companion. The latter system is a descendant of a high-mass X-ray binary (HMXB), in which the companion of the neutron star loses its mass through wind mass loss or a mass-transfer phase, exposing its helium core (e.g. Bhattacharya \\& van den Heuvel 1991). Although the existence of such a system is only found in Cyg X-3 (van Kerkwijk et al. 1992, 1996), a detailed study of the evolution of helium stars in binary systems with a compact companion is important as the systems form the bridge between the evolution of X-ray binaries and the formation of double compact-object binaries.  The evolution of a helium star in a binary system has been studied e.g. by Savonije, de Kool \\& van den Heuvel (1986) who evolved a 0.6~\\msun\\ non-degenerate helium star with a 1.3~\\msun\\ compact companion. With a short orbital period of $P = 37^{\\mathrm{m}}$, Roche-lobe overflow (RLOF) takes place during helium core burning. Ergma \\& Fedorova (1990) evolved helium stars with masses of 0.5, 0.766, and 1~\\msun. As the companion star, they took white dwarfs with the same range of masses in a combination such that the systems have a mass ratio of $0.5 \\leq M_{\\mathrm{He}}/M_{\\mathrm{WD}} \\leq 2$. The periods are also so short, of $26\\fm2$ -- $62\\fm6$, that the helium stars transfer mass to the companion while they are still burning helium in the centre. A similar work was also carried out by Tutukov \\& Fedorova (1990). A study of systems with more massive helium stars in wider orbits has been carried out by e.g. Delgado \\& Thomas (1981) who considered helium stars with masses of 2, 2.7, 3.3, and 4~\\msun\\ with a massive main-sequence companion. The helium star fills its Roche lobe after helium is exhausted in the core. Habets (1986a) evolved a 2.5~\\msun\\ helium star with a 17~\\msun\\ main-sequence companion in a wide orbit ($P = 20\\fd25$) such that mass-transfer phase occurs during first carbon burning convective shell.  In this work we study the evolution of helium stars with masses of 1.5 -- 6.7~\\msun\\ with a 1.4~\\msun\\ neutron-star companion. Such systems are assumed to be the remnants of Be/X-ray binaries with masses in the range of 8 -- 20~\\msun\\ which underwent mass transfer as case B (RLOF is initiated during hydrogen shell burning) or case C (during helium shell burning). As the result of the large mass ratio, mass transfer from the Be star to the neutron star is dynamically unstable and the two components are embedded in a common envelope (CE), leading to the spiraling-in of the neutron star in the envelope of the Be star. We assume that the system survived the common envelope and spiral-in process. We are interested to investigate the final fate of the systems (whether they will become white-dwarf/neutron-star or double neutron-star binaries), the type of supernovae (SN) they might produce, as well as to study which systems are stable to RLOF. A study on the similar range of mass (2 -- 6~\\msun\\ helium stars) has been carried out by Avila Reese (1993). With a Roche radius of 0.6 and 0.7~\\rsun, in that study RLOF takes place after helium core burning is terminated and the calculations were done up to the ignition of carbon in the centre. Apart from the slightly wider range of mass in our study compared to the latter work and a larger range of orbital periods used in our work, we also are able to follow the evolution of the helium stars to more advanced evolutionary stage, i.e. beyond the carbon ignition, which enables more detailed conclusions on the fate of the systems. We take 6.7~\\msun\\ as the upper limit of our calculation because more massive helium stars undergo a dynamically unstable mass transfer.  In Sect.~\\ref{helium:sec:evolution} we describe the possibility for the formation channel of a helium star in a binary system with a neutron-star companion, and some basic assumptions used in the calculation of the orbital evolution. The computational code, the input parameters used in the code, and the calculation that constrains the initial mass of the helium star and the initial period are described in Sect.~\\ref{helium:sec:method}. We discuss the results in two sections; Sect.~\\ref{helium:sec:BA} for the cases in which RLOF takes place during helium core burning and Sect.~\\ref{helium:sec:BB} for that during helium shell burning. Our conclusions are given in Sect.~\\ref{helium:sec:conclusion}. ",
        "conclusions": "\\label{helium:sec:conclusion}  \\begin{figure} \\centerline{\\psfig{file=mb825f11.eps,width=\\linewidth}} \\caption[]{The final fate of the helium stars borm in systems consisting of helium star and neutron star: CO white dwarf (solid-circle), ONe white dwarf (open-circle), or neutron star (solid-star). $P_{\\mathrm{i}}$ and $M_{\\mathrm{He}}$ are the orbital period and helium star mass at the onset of the evolution of the binary. The lines separate the regions of case BA, BB, and BC mass transfers, taken from Pols' (2002) single helium stars calculations. The shaded area mark the region where mass transfer is dynamically unstable.} \\label{helium:fig:final} \\end{figure}  We have done calculations of helium stars with masses in the interval of 1.5 -- 6.7~\\msun\\ with a 1.4~\\msun\\ neutron star companion. We allowed the helium star to fill its Roche lobe during helium core burning (case BA) and helium shell burning (case BB mass transfer). The final fate of the helium star as a function of its initial mass and period is summarized in Fig.~\\ref{helium:fig:final}.  Case BA mass transfer from helium stars with masses of 1.5 -- 2.9~\\msun\\ occurs during helium core burning and stops when the central abundance of helium drops below 0.1 and the star contracts. In helium stars less massive than 2.4~\\msun, RLOF is stable and takes place on the nuclear timescale. For $2.4 \\leq M_{\\mathrm{He}} / \\mathrm{M_{\\sun}} \\leq 2.9$, a rapid, thermal-timescale phase of RLOF is followed by a stable, slow phase of mass transfer on the nuclear timescale. Driven by radius expansion, a stable phase of case BAB mass transfer occurs during helium shell burning. This phase of mass transfer stops when the degeneracy border is crossed and the star contracts. During RLOF, the entire envelope is removed from the helium star. The remnants of case BA mass transfer are heavy, degenerate CO white dwarfs. This implies that a star with initial mass as large as 12~\\msun\\ still can become a white dwarf, as also found by e.g. Wellstein et al. (2001) in stars that go through conservative contact-free case A evolution.  Stable case BB mass transfer from 1.5 -- 2.1~\\msun\\ stars takes place on the nuclear timescale and stops when the degeneracy border is crossed and the star contracts. The entire envelope is removed during RLOF, and heavy, CO (from 1.5 -- 1.8~\\msun) or ONe (from 2.1~\\msun\\ helium stars) white dwarfs are the remnants in this range of mass. Case BB mass transfer from 2.4 -- 2.5~\\msun\\ helium stars produces ONe white dwarfs and that from more massive ones produces neutron stars. RLOF from $M_{\\mathrm{He}} \\geq 2.4$~\\msun\\ takes place on the thermal timescale. In close-orbit systems, this is followed by a slow, stable phase of mass transfer. There is a tendency that the helium envelope becomes convective in helium stars less massive than 4~\\msun. Mass transfer might become dynamically unstable in this case. However, the star probably collapses before a spiral-in can occur. Even if a CE and spiral-in phase does occur, the system is likely to survive. In both cases, a double neutron-star binary is formed.  From the amount of helium left in the envelope, we suggest that helium stars of high mass or in a wide orbit produce type Ib SNe. Type Ic SNe come from helium stars of lower mass or in a close orbit. We are also able to provide a zone of avoidance for the formation of double neutron stars."
    },
    "0201/astro-ph0201525_arXiv.txt": {
        "abstract": "We present the first results from an {\\it XMM-Newton} serendipitous medium-deep survey, which covers nearly three square degrees. We show the $log$ $N$ - $log$ $S$ distributions for the 0.5-2, 2-10 and 5-10 keV bands, which are found to be in good agreement with previous determinations and with the predictions of XRB models. In the soft band we detect a break at fluxes around $5\\times10^{-15}$ cgs. In the harder bands, we fill in the gap at intermediate fluxes between deeper {\\it Chandra} and {\\it XMM-Newton} observations and shallower {\\it BeppoSAX} and {\\it ASCA} surveys. Moreover, we present an analysis of the broad-band properties of the sources. ",
        "introduction": "While in the soft band (0.5-2 keV) {\\it ROSAT} (\\cite{abaldi-F:has98}) and especially {\\it Chandra} (\\cite{abaldi-F:ros01}) has resolved almost all the XRB, in the hard band (2-10 keV) the XRB has been resolved at a 25\\%-30\\% level with {\\it BeppoSAX} and {\\it ASCA} surveys (\\cite{abaldi-F:cag98}; \\cite{abaldi-F:gio00}) and recently at a 90\\% with {\\it Chandra} (\\cite{abaldi-F:ros01}). Moreover, in the very hard band (5-10 keV) the fraction resolved by {\\it BeppoSAX} is around 30\\% (\\cite{abaldi-F:fio99}) and recently in the {\\it XMM-Newton} Lockman Hole deep pointing a 60\\% is reached (\\cite{abaldi-F:has01}).\\\\ The optical counterparts of the objects making the XRB are predominantly Active Galactic Nuclei (AGN). In the soft band the predominant fraction is made by unabsorbed AGN, with a small fraction of absorbed AGN (\\cite{abaldi-F:sch98}). The fraction of absorbed type-2 AGN rises if we consider the spectroscopic identifications of hard X-ray sources in {\\it BeppoSAX}, {\\it ASCA} and {\\it Chandra} surveys (\\cite{abaldi-F:fio01}; \\cite{abaldi-F:del00}; \\cite{abaldi-F:toz01}).\\\\ The X-ray and optical observations are consistent with current XRB synthesis models (\\cite{abaldi-F:com95}; \\cite{abaldi-F:gil01}), which explain the hard XRB spectrum with an appropriate mixture of absorbed and unabsorbed AGN, by introducing the corresponding luminosity function and cosmological evolution. However, these models require the presence of a significant population of type-2 QSOs (\\cite{abaldi-F:nor01}), not yet detected in sufficient quantities. Type-2 QSOs are rare (so far, only a few are known), luminous and hard (heavily absorbed in the soft band). A good way of finding them is to perform surveys in the hard X-ray bands, covering large solid angles. The large throughput and effective area, particularly in the harder bands, make {\\it XMM-Newton} currently the best satellite to perform hard X-ray surveys.\\\\ In this poster contribution we present results from the HELLAS2XMM survey (\\cite{abaldi-F:bal02}), one of its main goals is to constrain the contribution of absorbed AGN to the XRB. ",
        "conclusions": "We are carrying out a serendipitous {\\it XMM-Newton} survey. We currently cover nearly three square degrees in 15 fields observed during satellite calibration and performance verification phase. This is, to date, the {\\it XMM-Newton} survey with the largest solid angle.\\\\ The main results can be summarized as follows: \\begin{itemize} \\item The $log$ $N$-$log$ $S$ relations in the 0.5-2 keV, 2-10 keV and 5-10 keV band are in agreement with previous determinations; \\item in the hard bands we sample an intermediate flux range: deeper than {\\it ASCA} and {\\it BeppoSAX} and shallower than {\\it Chandra} and {\\it XMM-Newton} deep surveys \\item We find an evidence for hard sources emerging below 0.5-2 keV fluxes of $2\\times10^{-14}$ erg cm$^{-2}$ s$^{-1}$ and 2-10 keV fluxes of 10$^{-13}$ erg cm$^{-2}$ s$^{-1}$. \\end{itemize}"
    },
    "0201/astro-ph0201243_arXiv.txt": {
        "abstract": "I review the current status of Fermi acceleration theory at relativistic shocks.  I first discuss the relativistic shock jump conditions, then describe the non-relativistic Fermi mechanism and the differences introduced by relativistic flows.  I present numerical calculations of the accelerated particle spectrum, and examine the maximum energy attainable by this process.  I briefly consider the minimum energy for Fermi acceleration, and a possible electron pre-acceleration mechanism. ",
        "introduction": "\\label{intro} A ubiquitous feature of astrophysical objects involving relativistic flows, such as active galactic nuclei (AGNs), gamma-ray bursts (GRBs) and Crab-like supernova remnants (SNRs), is the presence of nonthermal, power-law emission spectra (i.e. with flux density $F_\\nu \\propto \\nu^{-\\alpha}$, where $\\nu$ is the frequency and $\\alpha$ the spectral index), in particular in the radio and hard X-ray or gamma-ray domains. This emission is believed to be produced by accelerated particles having a corresponding power-law energy spectrum; more specifically, in most of these objects the emission is thought to be from accelerated electrons radiating via the synchrotron or inverse Compton mechanisms (see Mastichiadis, this volume).  The aim of the present review will be to discuss the probable mechanism of this acceleration and the spectra that may be expected theoretically.  The most widely invoked mechanism for the acceleration of particles to power-law spectra in non-relativistic contexts, such as SNR blast waves or interplanetary shocks, is \\emph{Fermi acceleration}.  It seems likely that shocks are responsible for particle acceleration in relativistic flows as well, and this is indeed explicitly assumed in models of GRBs and Crab-like SNRs.  It is then natural to consider how the Fermi mechanism could operate at relativistic shocks, and what the resulting spectrum would be.  The focus of this contribution will thus be the relativistic version of Fermi shock acceleration.  This review is organised as follows: in Sect.~\\ref{relshocks}, I discuss the shock jump conditions at relativistic shocks, emphasising the aspects relevant to particle acceleration; this section is intended to be self-contained.  In Sect.~% \\ref{Fermi}, I describe the Fermi acceleration mechanism in detail, first reviewing its main features in the context of non-relativistic shocks, and then presenting the resulting spectrum for ultra-relativistic and more moderately relativistic shocks. In Sect.~\\ref{maxmin}, I examine the acceleration time scale and the maximum energy attainable by this mechanism, and consider the minimum energy for Fermi acceleration of electrons in an electron--ion shock, and a possible pre-acceleration mechanism. ",
        "conclusions": "The shock velocity ratio $r$ across a relativistic shock is in general a function of the assumed upstream temperature as well as the shock Lorentz factor $\\Gammash$, but it rapidly tends to the ultra-relativistic limit $r=3$ for $\\Gammash \\gesim 10$.  The ultra-relativistic Fermi acceleration regime then mirrors some of the simplicity of the non-relativistic, strong shock regime, this asymptotic shock velocity ratio corresponding to an asymptotic power-law index of the accelerated particle distribution. For the specific case of isotropic direction-angle scattering on both sides of the shock, this spectral index is $p = 2.23 \\pm 0.01$; more generally, a value of $p$ in the range 2.2--2.3 is found under a variety of particle transport assumptions.  These values are consistent with the observed spectra of sources thought to contain ultra-relativistic shocks, such as gamma-ray burst afterglows and Crab-like supernova remnants.  For moderately relativistic shocks, the spectral index depends on the shock jump conditions as well as $\\Gammash$; in particular, shocks in a relativistic gas typically yield steeper spectral indices than the above ultra-relativistic values.  The maximum energy $E_\\mathrm{max}$ of the Fermi-accelerated particle distribution is determined by the acceleration time, which is in general set by the upstream residence time.  For acceleration at the unmodified, external blast wave of relativistic fireballs, this yields $E_\\mathrm{max} \\sim 10^{16}$\\,eV for typical parameters of the surrounding interstellar medium, ruling out the production of ultra-high-energy cosmic rays in this context.  If neutron star binary merger events give rise to relativistic blast waves with $\\Gammash \\gesim 10^3$, these can provide an alternative scenario for UHECR production: ions accelerated in the pulsar wind present before the merger can be boosted to energies $\\gesim 10^{20}$\\,eV by the blast wave with high efficiency; deceleration of the blast wave in the pulsar wind bubble yields a spectral index $p = 2$ and a typical lower cutoff around $3 \\times 10^{18}$\\,eV. There is also a minimum energy for Fermi acceleration, set by the requirement that the shock thickness be small relative to the particle Larmor radius.  In electron--ion shocks, this requires a distinct pre-acceleration mechanism for the electrons, which could be the resonant ion cyclotron wave acceleration mechanism of Hoshino et al.\\ \\cite{Hosetc92}."
    },
    "0201/astro-ph0201075_arXiv.txt": {
        "abstract": "The purpose of this article is to give a brief account of what we hope to learn from the future CMB experiments, essentially from the point of view of primordial cosmology. After recalling what we have already learnt,  the principles of parameter extraction from the data are summarized. The discussion is then devoted to the information we could gain about the early universe, in the framework of the inflationary scenario, or in more exotic scenarios like brane cosmology. ",
        "introduction": "I will start  by recalling briefly what we have already learnt from the Cosmic Microwave Background (CMB), which is already in itself rather impressive. For more details, the reader is invited to refer to the numerous reviews on the subject, for example the lecture notes of the 1993 les Houches school by R. Bond \\cite{bond96} and those of the 1999 les Houches school by F. Bouchet, J.L. Puget and J.M. Lamarre \\cite{bouchet99}.  As is well known, the CMB was predicted in 1948 \\cite{cmb48} and discovered less than twenty years later \\cite{cmb65}. Since then, it has been measured repeatedly, and with increasing precision. The CMB spectrum is  that of a black body to a very high precision. This feature  gives one of the strongest arguments in favour of the hot big bang model, according to which the photons where in thermal equilibrium in the past. This also implies some very stringent constraints on energy release in the universe after a redshift $z\\sim 10^{6-7}$.  A second important feature of the CMB is that it is almost isotropic but not quite. There is first a dipole at the $10^{-3}$ level, which is usually interpreted as the motion of Earth with respect to the CMB rest frame. There are then higher multipole anisotropies at the $10^{-5}$ level, which had been expected for a long time, and observed for the first time by the COBE satellite ten years ago \\cite{cobe92}. These anisotropies had been expected for a long time because they were believed (and are still today) to be generated at the moment of last scattering by very tiny cosmological fluctuations,  the ancestors of the present cosmological  structures.   Before going on, let us recall quickly the basic formalism to describe the CMB anisotropies. The temperature anisotropies can be expanded in terms of (scalar) spherical harmonics, \\beq {\\Delta T\\over T}(\\theta,\\phi)=\\sum_{l,m} a_{lm} Y_{lm}(\\theta, \\phi). \\label{alm} \\eeq For the theorist, the temperature is an isotropic random field and therefore the multipole coefficients $a_{lm}$ random variables. One can define the angular power spectrum by \\beq C_l=\\langle |a_{lm}|^2\\rangle, \\label{Cl} \\eeq which is enough to specify entirely  the temperature random field if it is Gaussian.  What is usually plotted is the quantity $l(l+1) C_l$. If one assumes coherent scale-invariant initial fluctuations, such as those produced by inflation,  one expects to observe a plateau at low $l$, corresponding to large angular scales. At smaller scales, one expects to see oscillations \\cite{oscillations}. The reason is that a given  Fourier mode, characterized by a constant comoving wavenumber $k$, will start to oscillate as soon as its wavelength is within the Hubble radius, i.e. when pressure enters into play. Of course, different Fourier modes enter the Hubble radius at different times and thus, at a given time, they are at different stages of their oscillatory pattern. The CMB, being then essentially a snapshot of the last scattering surface, we thus expect to see oscillations in the angular power spectrum.  By contrast, the   topological defect models, the main competitor facing  inflation for many years, do not predict oscillations because the contributions of many incoherent fluctuations, generated at different times by topological defects, add and smear out the oscillations \\cite{defects}.  With the recent data from Boomerang \\cite{boomerang}, Maxima \\cite{maxima} and DASI \\cite{dasi}, the picture which is now emerging is that predicted by inflation and not by topological defects, which cannot account for  the main part of the initial spectrum. Moreover  the position of the first peak suggests  that our Universe is quasi-flat.  In less than ten years, we have thus learnt a lot from the CMB anisotropies. The  coming decade should be extremely fruitful as well, with several planned experiments, the most ambitious being the Planck satellite mission \\cite{planck}. So far, it is however remarkable, for cosmology, that there has been no real surprise (apart a quantitative surprise with the amplitude of the anisotropies)  with the CMB. The simplest theoretical models were  able to  predict in advance what we have observed. The question is how long this situation is going to last. With the increasing precision of the forthcoming experiments, will the simplest early universe models survive or will one need slightly more complicated models, a lot of which have already been explored by theorists ? ",
        "conclusions": "It is always extremely difficult to make predictions about the future of a scientific domain and I will not take this risk. The amount of information we will learn from the future CMB observations will depend on how close or how far they turn out to be with respect to the {\\it canonical} version of the early universe model, that of an inflationary phase generated by a single field in slow-roll motion. To be too close or too far are probably the cases where the gain of information will be  minimal for theorists. Obviously, being very close would be  a tremendous success for the currently prefered model but would not bring any surprise. However, being too far might  not provide so much information as well, unless it corresponds to the predictions of a model already considered. We would have to revise some of our ideas but it is usually  easy for theorists to cook up a model, or even several, which will fit the data {\\it a posteriori}. The most stimulating situation for cosmology, where the number of observations concerning the early universe is extremely limited, would occur if  the new data roughly confirm the overall picture but add     details that reveal some deviations from the canonical model. This is the best situation to learn something because this is the case where one is most likely to interpret correctly the unexpected features of the data. Let us hope therefore than the next observations will put us in such a position.   \\ack It is a pleasure  to thank the organizers for this remarkable conference. I would also like to thank the {\\it Ambassade de France} in South Africa for their help and financial support."
    },
    "0201/astro-ph0201305_arXiv.txt": {
        "abstract": "{ We have carried out a quantitative trend analysis of the crystalline silicates observed in the ISO spectra of a sample of 14 stars with different evolutionary backgrounds. We have modeled the spectra using a simple dust radiative transfer model and have correlated the results with other known parameters. We confirm the abundance difference of the crystalline silicates in disk and in outflow sources, as found by Molster et al. (1999a). We found some evidence that the enstatite over forsterite abundance ratio differs, it is slightly higher in the outflow sources with respect to the disk sources. It is clear that more data is required to fully test this hypothesis. We show that the 69.0 micron feature, attributed to forsterite, may be a very suitable temperature indicator. We found that the enstatite is more abundant than forsterite in almost all sources. The temperature of the enstatite grains is about equal to that of the forsterite grains in the disk sources but slightly lower in the outflow sources. Crystalline silicates are on average colder than amorphous silicates. This may be due to the difference in Fe content of both materials. Finally we find an indication that the ratio of ortho to clino enstatite, which is about 1:1 in disk sources, shifts towards ortho enstatite in the high luminosity (outflow) sources.} ",
        "introduction": "The Infrared Space Observatory (ISO) has provided a new and unprecedented view on the occurence and composition of circumstellar and interstellar dust. One of the surprises of the ISO mission was the discovery of ubiquitous crystalline silicates in circumstellar dust shells of both evolved and young stars (see e.g. Waters et al. 1996; Waelkens et al. 1996). We have carried out an extensive study of the presence and properties of crystalline silicates. The present study is the third in a series, in which we study these silicates using ISO spectra. In previous papers (Molster et al. 2001c; 2001d; hereafter Papers I and II respectively) we have measured and described the circumstellar dust features found in the infrared spectra of 17 stars with different evolutionary status. The majority of these features could be identified with crystalline olivines (Mg$_{2x}$Fe$_{2-2x}$SiO$_4$) and pyroxenes (Mg$_{x}$Fe$_{1-x}$SiO$_3$), where $1 \\ge x \\ge 0$. J\\\"{a}ger et al. (1998, hereafter JMD) measured the mass absorption coefficient of crystalline pyroxenes and olivines  with different Fe over Mg ratios. Bands of both materials show a shift in the wavelength position of the peaks to longer wavelengths with increasing Fe content. The detection of the 69 micron feature, which is very sensitive to the Fe/Mg ratio (Koike et al. 1993; JMD), as well as the relative strength of the crystalline silicate features in the spectrum of IRAS09425-6040 (Molster et al. 1999a), led to the conclusion that the crystalline olivines observed in the ISO spectra are very Mg-rich ($x > 0.95$); the Mg-rich end member of the olivines is called forsterite. Similarly, the enstatite band at 40.5$\\mu$m is sensitive to the Fe/(Fe+Mg) ratio and points top the presence of Mg-rich pyroxenes. The identification of the dust species is very important for a better insight in the formation and evolution of dust. This may lead to a better understanding of the mass loss process and thus the evolution of the mass-losing star itself.  There is a clear separation between sources with and without a dusty disk. This difference is evident quantitatively in the sense that the crystalline silicate features are stronger with respect to the continuum in the sources which are surrounded by a disk (Molster et al. 1999a), and also qualitatively in the shape of the features, which is a proof for different dust properties (Paper I).  However, more quantitative statements are necessary to come to a better understanding of the nature of the circumstellar dust in these objects, and of their formation and processing history. In order to get these quantitative statements, a comparison with laboratory measurements is necessary. Unfortunately the laboratory measurements do not always agree with each other. In Paper II we discussed different laboratory measurements of olivines and pyroxenes, and possible causes of discrepancy. Despite these differences, qualitative agreement with the ISO spectra is already quite impressive as we will demonstrate in the present study.  In Section~\\ref{sec:trends} we discuss trends in the position and width of the solid state bands. In Section~\\ref{sec9:model} we apply a very simple optically thin dust model to the spectra to determine a typical temperature for the dust species. The results of this modelling are used to look for correlations which are discussed in Section~\\ref{sec:corr}.  Peak positions show variation from source to source. We adopt the naming in Paper I, which implies that if we refer to a wavelength position we use $\\mu$m, while if we refer to the name of specific feature we will write `micron'. ",
        "conclusions": "We can summarize the main results of this study as follows: \\begin{itemize} \\item[1] The ISO spectra could be reasonably well fitted with laboratory spectra of forsterite and enstatite. \\item[2] The models underestimate the flux at 18, 29.6, 30.6, 48 $\\mu$m and sometimes at 40.5 $\\mu$m, which is an indication for the presence of (an)other dust component(s). No convincing identification could be made yet. Diopside does have features at most of these wavelengths, but also strong features at others which are weak or absent in the ISO-spectra. \\item[3] The 19.5 micron feature is often overestimated by our model spectra. No explanation is yet known for this phenomenon, but it should be noted that in the full radiative transfer modelling it appeared to be much less of a problem. This might indicate that optical depth effects play a role. Also the calculation of the absorption coefficients from the constants instead of the absorption coefficients from laboratory particles might lead to differences. \\item[4] The band width of the laboratory data is larger than in our ISO spectra. This difference is likely a temperature effect, and might be used as an independent temperature indicator. Especially the 69.0 micron band is very suitable for this analysis. \\item[5] The temperature of the forsterite and enstatite grains are (almost) similar for the disk sources, while the forsterite is slightly warmer in the outflow sources. This would imply that the forsterite and/or enstatite grains differ slightly in the disk and outflow sources. It is not clear whether this difference is due to a different formation process, or due to a different dust process history after the grain formation. Since this trend is only based on 3 sources more data is required to confirm the difference between the dust and outflow sources. \\item[6] The crystalline silicates are colder than the amorphous silicates. This is probably due to the difference in chemical composition. No Fe is present in the crystalline silicates, while in the amorphous silicates it is expected to explain the higher absorptivity. This difference in temperature also implies that the crystalline and amorphous grains are two distinct grain populations. \\item[7] Enstatite is on average a factor 3--4 more abundant than forsterite in our sources. There are indications that the enstatite over forsterite ratio in the outflow sources is higher than in the disk sources. \\item[8] In the low luminosity sources the spectra were well fitted with an equal amount of ortho- and clino-enstatite, while in the two high luminosity sources more ortho-enstatite seems to be present. \\item[9] No correlation could be found between the crystallinity and the temperature of the dust. Also the luminosity of the stars does not seem to be correlated with the enstatite over forsterite ratio. \\item[10] These simple model fits already give a good insight in the dust around our stars. In Paper I the shape of the features naturally separated the disk and outflow sources. In this study the differences between these two categories became again evident. \\end{itemize}  \\vspace{1.0cm} \\noindent{\\bf Acknowledgements.}\\\\ We greatly thank Janet Bowey for providing her laboratory data. FJM wants to acknowledge the support from NWO under grant 781-71-052 and under the Talent fellowship programm. LBFMW acknowledges financial support from an NWO `Pionier' grant."
    },
    "0201/astro-ph0201133_arXiv.txt": {
        "abstract": "We search the BOOMERanG maps of the anisotropy of the Cosmic Microwave Background (CMB) for deviations from gaussianity. In this paper we focus on analysis techniques in pixel-space, and compute skewness, kurtosis and Minkowski functionals for the BOOMERanG maps and for gaussian simulations of the CMB sky. We do not find any significant deviation from gaussianity in the high galactic latitude section of the 150 GHz map. We do find deviations from gaussianity at lower latitudes and at 410 GHz, and we ascribe them to Galactic dust contamination. Using non-gaussian simulations of instrumental systematic effects, of foregrounds, and of sample non-gaussian cosmological models, we set upper limits to the non-gaussian component of the temperature field in the BOOMERanG maps. For fluctuations distributed as a 1 DOF $\\chi^2$ mixed to the main gaussian component our upper limits are in the few \\% range. ",
        "introduction": "\\label{sec:intro}  In most inflationary scenarios the primordial density field is expected to be gaussian \\citep[]{Pee99}. In contrast, structure formation scenarios based on topological defects \\citep[]{Avel98} or less general inflationary models \\citep[]{Linde97, Mart00, Cont99} predict a non-gaussian density field. Thus, measurement of the statistical nature of the CMB anisotropies can distinguish between these scenarios \\citep[]{Fisc85}.  The presence of noise in the measurements combined with the limited coverage of the present observations can mask cosmological non-gaussian features. Moreover, the presence of systematic effects can produce subtle instrumental non-gaussian features in intrinsically gaussian anisotropy maps.  Efforts to identify non-gaussianities in the CMB have been extensively carried out for the COBE data \\citep[]{Fer98, Band99, Kom01, Brom99} and no significant detection of cosmological non-gaussianity has been reported \\citep{MagZZ}. However, the sensitivity of COBE-DMR to the expected levels of cosmological non-gaussianity from rare highly non-linear events like topological defects \\citep[]{Durr96} is not very high, due to the large field of view which smears out the effects of small scale features. While the present observations on the power spectrum are inconsistent with a structure formation scenario solely based on defects \\citep[]{Durr98}, a mixed inflation+defects model is still compatible with the data \\citep[]{Bouc01}.  Furthermore, the low signal-to-noise ratio and coarse resolution in the COBE maps also makes it difficult to detect primordial non-gaussianity; the central-limit theorem states that the sum of various instrumental effects will make the distribution tend to a gaussian \\citep[]{Novi00}.   Analyses to date of higher angular resolution experiments, like QMASK \\citep[]{Park01} and MAXIMA \\citep[] {Sant01, Wu01}, show full consistency with gaussianity.  In this paper we focus on the maps produced by the BOOMERanG experiment \\citep[]{deBe2000,Nett2001}. Due to its wide sky and frequency coverage, BOOMERanG is ideally suited to carry out an accurate analysis of the possible systematic effects present in the detected signal. We analyze them using a Monte-Carlo approach, in order to set quantitative upper limits for the level of primordial non-gaussian fluctuations present in the maps.  The data we use are presented in section ~\\ref{sec:data}.  We use five pixel-space estimators of departures from gaussianity: the skewness and kurtosis of the CMB temperature distribution, and the three Minkowski functionals: area, length and genus \\citep[]{Mink03}. In section ~\\ref{sec:Analy} we apply these estimators to the BOOMERanG maps and to gaussian Monte-Carlo simulations of the CMB sky observed by BOOMERanG. In section ~\\ref{sec:Syst} we analyze template maps of other non-gaussian signals which could in principle be present in the BOOMERanG maps; we study how sensitive the pixel-space techniques are in detecting these effects, and we estimate how much these affect the measurements of cosmological non-gaussian signals. In section ~\\ref{sec:sour} we compare our measurements to the cosmological expectations for a few sample scenarios.  The methods presented here are especially useful in the detection of highly non-linear features (as expected in topological defects theories or for Galactic foregrounds) \\citep[]{Phil01}. One may alternatively use a wavelets approach \\citep[]{Hiv01}, or a bispectrum approach \\citep[]{Cont01}. ",
        "conclusions": "\\label{sec:Disc}  The pixel-based techniques discussed here and applied to the central region observed by BOOMERanG at 150 GHz confirm the Gaussianity of the detected CMB fluctuations, and exclude with high confidence $\\chi^2$ distributed CMB temperature fluctuations. These techniques also detect non-Gaussian fluctuations due to interstellar dust at high frequencies and at lower latitudes. Subdominant non-Gaussian fluctuations mixed to the Gaussian ones are strongly excluded for some models, while are not detected at all in other cases."
    },
    "0201/astro-ph0201419_arXiv.txt": {
        "abstract": "We have begun a systematic numerical study of the nonlinear growth of nonaxisymmetric perturbations during the ambipolar diffusion-driven evolution of initially magnetically subcritical molecular clouds, with an eye on the formation of binaries, multiple stellar systems and small clusters. In this initial study, we focus on the $m=2$ (or bar) mode, which is shown to be unstable during the dynamic collapse phase of cloud evolution after the central region has become magnetically supercritical. We find that, despite the presence of a strong magnetic field, the bar can grow fast enough that for a modest initial perturbation (at $5\\%$ level) a large aspect ratio is obtained during the isothermal phase of cloud collapse. The highly elongated bar is expected to fragment into small pieces during the subsequent adiabatic phase. Our calculations suggest that the strong magnetic fields observed in some star-forming clouds and envisioned in the standard picture of single star formation do not necessarily suppress bar growth and fragmentation; on the contrary, they may actually promote these processes, by allowing the clouds to have more than one (thermal) Jeans mass to begin with without collapsing promptly. Nonlinear growth of the bar mode in a direction perpendicular to the magnetic field, coupled with flattening along field lines, leads to the formation of supercritical cores that are triaxial in general. It removes a longstanding objection to the standard scenario of isolated star formation involving subcritical magnetic field and ambipolar diffusion based on the likely prolate shape inferred for dense cores. Continuted growth of the bar mode in already elongated starless cores, such as L1544, may lead to future binary and multiple star formation. ",
        "introduction": "\\label{sec:introduction}  A basic framework has been developed for the formation of low-mass stars in relative isolation \\citep{FShu87}. It involves gradual condensation of dense cores from strongly magnetized, background molecular clouds through ambipolar diffusion, followed by dynamic core collapse to form stars. Quantitative studies based on this by now ``standard'' scenario have been carried out by many authors, with increasingly sophisticated input physics \\citep{TMouschovias99}. However, with few exceptions \\citep{RInde00, ABoss00}, axisymmetry has been adopted. The adopted symmetry precludes a detailed investigation of cloud fragmentation, generally thought to be a necessary step in the formation of binary and multiple stellar systems; it is in such systems that most stars are found. We seek to improve this situation by removing the restriction of axisymmetry and systematically investigating the effects of the strong magnetic fields envisioned in the standard picture of single star formation on fragmentation and their implications on binary and multiple star formation.  Fragmentation of nonmagnetic clouds has been studied extensively over the years, mostly through numerical simulations \\citep[and references therein]{PBodenheimer00}. In the canonical case of a spherical cloud, the fragmentation is mainly controlled by the ratios of the cloud thermal and rotational energy to the gravitational energy, $\\alpha$ and $\\beta$, although the distributions of density and angular momentum also play a role. Molecular line observations \\citep{AGoodman93} suggest that star-forming cores of molecular clouds are far from being rotationally supported, with a typical (low) value of $\\beta \\sim 0.02$. For such slowly rotating cores, the criterion for fragmentation is roughly $\\alpha \\lesssim 0.5$, if the cores are idealized as spheres of uniform density and rigid rotation \\citep{TTsuribe99}. More realistic, centrally peaked density distributions would lower the critical $\\alpha$ for fragmentation. Such distributions are commonly inferred in the well-observed starless cores such as L1544 \\citep{PAndre00, NEvans01}. The central concentration, if taken as an initial condition, tends to make fragmentation more difficult \\citep{ABoss93}.   Fragmentation of strongly magnetized clouds with mass-to-flux ratios below the critical value $(2\\pi)^{-1}G^{-1/2}$ (i.e., subcritical) is expected to be quite different \\citep{DGalli01}. Through magnetic braking, the field controls the angular momentum evolution of the cloud, until a supercritical dense core forms \\citep{SBasu94}; it thus determines the value of $\\beta$ of the core. More importantly, the field can provide most of the cloud support against self-gravity and thus allow a cloud to have an arbitrarily small thermal energy compared to the gravitational energy and still be in a mechanical equilibrium initially. The low value of $\\alpha$ is a key ingredient of fragmentation. On the other hand, the presence of a strong magnetic field tends to stifle cloud fragmentation, if the field and matter are well coupled \\citep{GPhillips86}. Indeed, for a frozen-in subcritical field, fragmentation is suppressed altogether \\citep{TNakano88}. Fragmentation can in principle resume in the part of cloud that has become magnetically supercritical \\citep{ABoss00}, through ambipolar diffusion. It is however difficult to predict {\\it a priori} how the fragmentation should proceed, if at all, in a cloud that is only {\\it partially} coupled to a strong magnetic field.  \\citet{WLanger78} showed through linear analysis that, in a lightly ionized medium such as a molecular cloud, the Jeans instability is not completely inhibited no matter how strong the magnetic field is; rather, it grows on an ambipolar diffusion, rather than dynamic, timescale. \\citet{ZLi01} followed the nonlinear evolution of a set of magnetically subcritical, super-Jeans mass clouds assuming axisymmetry and found that either a dense, supercritical core or ring forms as a result of ambipolar diffusion. The supercritical ring is expected to fragment readily into a number of dense cores, producing perhaps a small group of stars. The ring breakup will be studied elsewhere. In this Letter, we shall concentrate on the more subtle problem of the fragmentation of magnetized single cores, with an eye on forming binaries and multiple stellar systems. As is common with numerical studies of cloud fragmentation and binary formation, we focus on the growth of $m=2$ (or bar) mode \\citep[e.g.,][]{FNakamura97, PBodenheimer00}. In \\S~\\ref{sec:models} we describe our formulation of the problem of nonaxisymmetric cloud evolution. Numerical results on bar formation are presented in \\S~\\ref{sec:results}, and their implications on binary and multiple star formation as well as the observed elongation of dense cores are discussed in \\S~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  We have followed numerically the growth of $m=2$ mode during the ambipolar-diffusion driven evolution of magnetically subcritical clouds. Our main conclusion is that, despite the presence of the strong magnetic field, a perturbation of modest amplitude can grow nonlinearly into a highly elongated bar, which is expected to fragment into small pieces gravitationally. Before discussing fragmentation, we comment on the effects of bar growth on the observed shapes of molecular cloud cores.  Dense cores of molecular clouds are intimately associated with star formation, with roughly half of them already harboring infrared sources \\citep{CBeichman86}. The other half are thought to be well on their way to star formation, with the ``starless'' phase lasting for only a few dynamic times. The cores are observed to have significant elongation, with a typical aspect ratio of 2 \\citep{PMyers91}. Statistics of core elongation have been interpreted as indicating most cores have prolate 3-dimensional shape \\citep{BRyden96, CCurry01}, although oblate cores formed as a result of settling along magnetic field lines can have a projected aspect ratio similar to that observed \\citep{ZLi96}. Nonlinear growth of the bar mode in a direction perpendicular to the field lines modifies the core shape drastically, making it triaxial in general. Some evidence for the triaxial nature of cores has been marshalled by \\citet{SBasu00}. We propose that it is due to the significant bar growth during the transition period when $\\Gamma$ decreases from $\\sim 1$ to $\\sim 0.5$, after the core has become supercritical (which makes bar growth possible) but before a rapid collapse sets in (which leaves little time for the bar to grow). The bar growth removes an oft-invoked objection against the standard scenario of ambipolar diffusion-driven core formation based on the elongated or filamentary shape of dense cores \\citep{DWard-Thompson99}. Indeed, the triaxial nature of the dense cores, coupled with the relatively slow, subsonic infall motion inferred  \\citep{CLee01}, may provide the strongest support yet for the standard scenario.  The continued evolution of bar-like cores, such as the one shown in Figure \\ref{fig:1}c, would formally lead to a singular filament, if the isothermal assumption is kept. However, it is well known that the equation of state stiffens when the optical depth exceeds unity \\citep[see however][]{HMasunaga99}. The stiffening slows down the collapse, allowing more time for the elongated bar to fragment gravitationally into small pieces. In fact, the major axis of the bar exceeds twice the critical wavelength for fragmentation of an infinitely-long filament \\citep{RLarson85} in panels (e) and (f) of Figure \\ref{fig:1}. Such bar formation is also shown by \\citet{TMatsumoto99} and \\citet{LSigalotti01} for nonmagnetized clouds. Recent examples of (nonmagnetic) bar fragmentation induced by the stiffening of equation of state are given in \\citet{ABoss00b} and \\citet{LSigalotti01}. For magnetized bars, one needs to consider in addition the rapid decoupling of magnetic field from matter, which occurs above a density of order $10^{10}$~cm$^{-3}$ \\citep{RNishi91}. The decoupling decreases the magnetic support quickly, and is in some sense equivalent to a sudden cooling. It should make the fragmentation easier.   \\citet{ABoss00} studied the fragmentation of 3D magnetic clouds numerically, treating the magnetic forces and ambipolar diffusion in an approximate way. He concluded that magnetic fields can enhance cloud fragmentation, by reducing the tendency for the development of a central singularity, which would make fragmentation more difficult. We also find that magnetic fields can promote fragmentation, but for a different reason. Strong magnetic fields can support clouds of more than one Jeans mass, which provides an initial condition that is conducive to fragmentation once the magnetic support weakens. We believe that it is the multi-Jeans mass nature of the magnetically subcritical clouds that drives the nonlinear growth of the bar (and higher order) mode. We conclude that the magnetically subcritical clouds envisoned in the standard scenario of star formation can produce not only single stars but also, perhaps even preferentially, binary and multiple star systems. A parameter survey will be carried out to firm up this conclusion."
    },
    "0201/gr-qc0201069_arXiv.txt": {
        "abstract": "The endpoint of black hole evaporation is a very intriguing problem of modern   physics.   Based    on   Einstein-dilaton-Gauss-Bonnet   four dimensional  string gravity  model  we show that  black  holes do  not disappear  and should  become  relics at the  end  of the  evaporation process. The possibility of experimental  detection  of  such  remnant black holes is investigated. If they really exist, these objects could be  a  considerable  part  of  the non  baryonic  dark  matter  in our Universe. ",
        "introduction": "Theoretical physics faces nowadays a  great  challenge.  There is four dimensional Standard Model  on one side (and the additional dimensions are  not  required   to   explain  experimental  data)  together  with inflationary cosmology  based  on additional scalar fields \\cite{c01}. On  the  other  side  there   is   the   completely   supersymmetrical string/M-theory. Building links between those approaches \\cite{c02} is a very motivating goal of  modern  physics which could be achieved  by the study of microscopic black holes.  As  General  Relativity  is  not renormalizable, its  direct  standard quantization is  impossible.  To  build  a semiclassical gravitational theory, the usual  Lagrangian should be generalized, which is possible in different  ways. One of them is  to study  the action expansion  in scalar  curvature,  {\\it i.e.} higher order curvature corrections.  At the level of second order, according to the perturbational approach of string theory,  the most natural  choice is the 4D curvature invariant Gauss-Bonnet term $S_{GB} =  R_{ijkl}R^{ijkl}  - 4 R_{ij}R^{ij} + R^2$ \\cite{c03}.  With 4D action, it is not possible to consider only  $S_{GB}$ because, being full derivative, it does not contribute to  the field equations. It  must be  connected  it  with  a scalar  field $\\phi$  to  make its contribution dynamical. The following 4D effective  action with second order curvature corrections can be built: \\begin{eqnarray*} S & = & \\int d^4 x \\sqrt{-g} \\Biggl[ - R + 2 \\partial_\\mu \\phi \\partial^\\mu \\phi + \\lambda \\xi(\\phi) S_{GB} + \\ldots \\Biggr], \\end{eqnarray*} where $\\lambda$ is the string coupling constant. As  in cosmology, the most simple generalization of  the  theory (a single additional scalar field)  is  not  possible  because  while   dealing  with  spherically symmetric solutions, the ``no-hair'' theorem restriction must be taken into account.  Treating   $\\phi$   as  a  dilatonic  field,  the  coupling   function $\\xi(\\phi)$ is  fixed from the first  string principles and  should be written $\\exp(-2\\phi)$ \\cite{c04,c05}, which leads to : \\begin{eqnarray}\\label{a11} S & = & \\int d^4 x \\sqrt{-g} \\Biggl[ - R + 2 \\partial_\\mu \\phi \\partial^\\mu \\phi + \\lambda e^{-2\\phi} S_{GB} + \\ldots \\Biggr]. \\end{eqnarray} Such type of actions  can be considered as one of the  possible middle steps between General Relativity  and  Quantum Gravity. In this paper, we show that this effective string gravity model and its solutions can be applied for a description  of  the last stages of primordial  black holes (PBH)  evaporation  \\cite{c071,c072}  and suggests possible dark matter candidates \\cite{c08}. This should be understood in the general framework of Gauss-Bonnet black hole (BH) theory \\cite{c06,c07}.  The paper is  organized  as follows: In Section  II  we briefly recall previously obtained results  and point out some new features important for this study, Section III is devoted to the establishment of the new Hawking  evaporation  law (especially for the detailed description  of last stages of  Gauss-Bonnet BH evaporation),  in Section IV  we  show that the direct  experimental registration of such PBHs is impossible, Section V is  devoted  to  PBH relics as dark  matter  candidates  and Section VI contains discussions and conclusions. ",
        "conclusions": "In this paper,  the  BH type solution of  4D  effective string gravity action  with  higher order curvature corrections was  applied  to  the description of  BHRs. A corrected  version of the evaporation law near the minimal BH mass was  established.  It was shown that the  standard Bekenstein-Hawking  evaporation   forluma  must  be  modified  in  the neighbourhood of the last stages. Our main conclusion is to  show that contrarily to what is usually thought the evaporation does not  end up by the emission of  a few quanta with energy around Planck  values but goes  asymptotically  to  zero  with an infinite  characteristic  time scale.  The direct experimental registration of the products of evaporation of BHRs is  impossible. This gives  an opportunity to consider these BHRs as one of the main candidates for cold dark matter in our Universe."
    },
    "0201/astro-ph0201504_arXiv.txt": {
        "abstract": "}[1]{{ \\footnotesize \\noindent {\\bf Abstract} #1 \\\\}}  \\renewcommand{\\author}[1]{\\subsubsection*{#1}} \\newcommand{\\address}[1]{\\subsubsection*{\\it#1}}  \\setlength{\\textheight}{20cm} \\setlength{\\textwidth}{13.5cm}   \\newcommand{\\etal}{et al.\\ } \\newcommand{\\1}{\\ensuremath{^{-1}}} \\newcommand{\\2}{\\ensuremath{^{-2}}} \\newcommand{\\3}{\\ensuremath{^{-3}}} \\newcommand{\\PC}{\\ensuremath{\\mathrm{\\,pc}}} \\newcommand{\\KPC}{\\ensuremath{\\mathrm{\\,kpc}}} \\newcommand{\\MPC}{\\ensuremath{\\mathrm{\\,Mpc}}} \\newcommand{\\KM}{\\ensuremath{\\mathrm{\\,km}}} \\newcommand{\\CM}{\\ensuremath{\\mathrm{\\,cm}}} \\newcommand{\\SEC}{\\ensuremath{\\mathrm{\\,s}}} \\newcommand{\\YR}{\\ensuremath{\\mathrm{\\,yr}}} \\newcommand{\\ERG}{\\ensuremath{\\mathrm{\\,erg}}} \\newcommand{\\ARCSEC}{\\ensuremath{\\mathrm{\\,arcsec}}} \\newcommand{\\MAG}{\\ensuremath{\\mathrm{\\,mag}}} \\newcommand{\\Msun}{\\ensuremath{\\,\\mathrm{M}_\\odot}} \\newcommand{\\Rsun}{\\ensuremath{\\,\\mathrm{R}_\\odot}} \\newcommand{\\Lsun}{\\ensuremath{\\,\\mathrm{L}_\\odot}} \\newcommand{\\MB}{\\ensuremath{M_\\mathrm{B}}} \\newcommand{\\Kp}{\\ensuremath{K^\\prime}} \\newcommand{\\Hubble}{\\ensuremath{\\mathrm{H}_\\circ}} \\newcommand{\\xten}[1]{\\ensuremath{\\times 10^{#1}}} \\newcommand{\\ten}[1]{\\ensuremath{10^{#1}}} \\newcommand{\\farcs}{\\ensuremath{.\\!\"}} \\newcommand{\\parfrac}[2]{\\ensuremath{\\left( \\frac{#1}{#2} \\right)}} \\newcommand{\\kms}{\\KM\\SEC\\1} \\newcommand{\\BH}{\\ensuremath{\\mathrm{BH}}} \\newcommand{\\BHs}{\\ensuremath{\\mathrm{BHs}}} \\newcommand{\\Mdot}{\\ensuremath{\\dot{M}}} \\newcommand{\\MBH}{\\ensuremath{M_\\mathrm{BH}}} \\newcommand{\\RBLR}{\\ensuremath{R_\\mathrm{BLR}}} \\newcommand{\\LEdd}{\\ensuremath{L_\\mathrm{Edd}}} \\newcommand{\\Lbol}{\\ensuremath{L_\\mathrm{bol}}} \\newcommand{\\Mbulge}{\\ensuremath{M_\\mathrm{Bulge}}} \\newcommand{\\Lbulge}{\\ensuremath{L_\\mathrm{Bulge}}} \\newcommand{\\rhobh}{\\ensuremath{\\rho_\\mathrm{BH}}} \\newcommand{\\rbh}{\\ensuremath{r_\\mathrm{BH}}} \\newcommand{\\thbh}{\\ensuremath{\\theta_\\mathrm{BH}}} \\newcommand{\\sigmastar}{\\ensuremath{\\sigma_\\star}} \\newcommand{\\mlr}{\\ensuremath{\\Upsilon}} \\newcommand{\\chisq}{\\ensuremath{\\chi^2}} \\newcommand{\\Sgra}{\\ensuremath{\\mathrm{SgrA}^\\star}} \\newcommand{\\aap}{A\\&A} \\newcommand{\\aj}{AJ} \\newcommand{\\apj}{ApJ} \\newcommand{\\apjl}{ApJ} \\newcommand{\\araa}{ARA\\&A} \\newcommand{\\mnras}{MNRAS} \\newcommand{\\nat}{Nature} \\newcommand{\\pasp}{PASP}    \\begin{document}   \\chapter*{Black Holes in Galactic Nuclei: the Promise and the Facts}  \\author{Alessandro Marconi}   \\address{Osservatorio Astrofisico di Arcetri, Largo E. Fermi 5, I-50125 Firenze, Italy}   \\abstract{It has long been suspected that Active Galactic Nuclei are powered by accretion of matter onto massive black holes and this belief implies their presence in the nuclei of most nearby galaxies as \"relics\" of past activity. Just a few years ago this was considered a paradigm but, recently, new ground-based and Hubble Space Telescope observations are producing a breakthrough in our knowledge on massive black holes.  I will review the evidence for the existence of black holes in galactic nuclei and how their presence is related to host galaxy properties and AGN activity.} ",
        "introduction": "In November 1783 John Michell presented to the Royal Society his idea of a {\\em dark star}, a star so massive that the escape velocity from its surface is larger than the speed of light. Combining the corpuscolary theory of light with Newton's theory of gravitation, he found that a star with the same density as the Sun but escape velocity equal to $c$ would have radius $R=486\\Rsun$ and mass $M=1.2\\xten{8}\\Msun$.  Michell also pointed out that, although dark stars are invisible, their presence could be inferred from the motion of other luminous bodies orbiting around them.  Similar ideas were independently presented in 1796 by Laplace in his {\\it \"Exposition du syst\\`eme du monde\"}.  In 1916 Karl Schwarzschild presented his exact solution of Einstein field equations deriving the well-known {\\em Schwarzschild radius} which Michell had {\\it exactly} determined although starting from {\\it wrong} assumptions.  The term {\\em black hole}, which is now commonly used, was not coined until 1967 by John Wheeler and the first observational evidence for the existence of a black hole was given in the early 70s by the observations of the binary X-ray source Cygnus X-1.  The discovery of quasars with their enormous energy output from small volumes of space suggested that they were powered by accretion of matter onto very massive black holes residing in galactic nuclei \\cite{lynden}.  The first observational evidence was found in the galaxy M87 whose nucleus seemed to host a 5\\xten{9}\\Msun\\ black hole \\cite{sarg}. The supermassive black holes (hereafter \\BHs) hosted in galactic nuclei, with masses in the range $\\ten{6}-\\ten{10}\\Msun$, are the topic of this review. Observational evidences for the existence of \\BHs\\ in galactic nuclei up to the early '90s are summarized in a review by Kormendy \\& Richstone \\cite{kr}.  At that time only a handful of \\BHs\\ were known from ground-based and early HST observations. ",
        "conclusions": ""
    },
    "0201/astro-ph0201218_arXiv.txt": {
        "abstract": "A stellar wind module has been developed for the \\phoenix\\ stellar atmosphere code for the purpose of computing non-LTE, line-blanketed, expanding atmospheric structures and detailed synthetic spectra of hot luminous stars with winds.  We apply the code to observations of Deneb, for which we report the first positive detections of mm and cm emission (obtained using the SCUBA and the VLA), as well a strong upper limit on the 850$\\mu$m flux (using the HHT).  The slope of the radio spectrum shows that the stellar wind is partially ionized.  We report a uniform-disk angular diameter measurement, $\\overline{\\theta_{\\rm UD}} = 2.40 \\pm 0.06\\ {\\rm mas}$, from the Navy Prototype Optical Interferometer (NPOI).  The measured bolometric flux and corrected NPOI angular diameter yield an effective temperature of $8600 \\pm 500\\ {\\rm K}$.  Least-squares comparisons of synthetic spectral energy distributions from 1220 \\AA\\ to 3.6 cm with the observations provide estimates for the effective temperature and the mass-loss rate of $\\simeq 8400 \\pm 100$ K and $8 \\pm 3 \\times 10^{-7}$ \\Mspyr, respectively. This range of mass-loss rates is consistent with that derived from high dispersion UV spectra when non-LTE metal-line blanketing is considered.  We are unable achieve a reasonable fit to a typical \\ha\\ P-Cygni profile with any model parameters over a reasonable range.  This is troubling because the \\ha\\ profile is the observational basis for Wind Momentum-Luminosity Relationship. ",
        "introduction": "\\label{section_a1}  A-type supergiants are the brightest stars at visual wavelengths (up to $M_V \\simeq -9$) and are therefore among the brightest single stars visible in galaxies.  For this reason, these stars have been of increasing interest in extragalactic astronomy where they show potential as independent distance indicators \\cite[][hereafter K99]{bresolin2001,kud99}.  This potential lies in the use of the Wind Momentum-Luminosity Relationship (WMLR) which is derived from Sobolev radiation-driven stellar wind theory (K99 and references therein). Testing this relationship and the theory of radiation driven winds most critically requires an accurate determination of the stellar mass-loss rates.  The WMLR states that $\\dot{M}v_{\\infty} \\propto L^{1/\\alpha}/R_\\star^{0.5}$, where \\mdot\\ is the mass-loss rate, \\vterm\\ is the terminal velocity of the wind, $R_\\star$ is a photospheric reference radius, and $\\alpha$ is a parameter related to the distribution of atomic line strengths for the spectral lines which drive the wind.  The application of the WMLR to nearby galaxies requires a calibration of the relationship using similar stars in the Milky Way.  To determine mass-loss rates for A-type supergiants, K99 modeled the hydrogen Balmer lines \\ha\\ and \\hg\\ for six stars in the Galaxy and M31 using a unified stellar atmosphere model. In their analysis, K99 synthesized the observed line profiles by adjusting model values for the mass-loss rate, the velocity law exponent $\\beta$, and a line-broadening parameter, $v_t$, after first adopting parameters for the effective temperature, \\Teff, the radius, $R_\\star$, the gravity, \\Logg, and the projected stellar rotational velocity.  The use of the \\ha\\ profile as the sole diagnostic of A-type supergiant wind momenta is problematic because \\ha\\ emission observed in stars of this type is known to be strongly variable \\cite[]{rose72,kaufer96}.  Nevertheless, the accuracy of mass-loss rate determinations for A-type supergiants in M31 is claimed to be at the level of 15\\% \\cite[]{mccarthy97}.  However, there has been no consensus on the mass-loss rate for Deneb ($\\alpha$~Cyg, HD~197345, spectral type A2Ia according to \\cite{mk73}), the brighest, nearest, and best-studied A-type supergiant, which we might expect to serve as an archetype and benchmark.  Published estimates of its mass-loss rate range over more than three orders of magnitude, from methods which include modeling the line profiles \\ion{Mg}{2} $\\lambda$2802 \\cite[][$3.1\\times10^{-9} - 1.5\\times10^{-7}$ \\Mspyr]{kp81} and \\ha\\ \\cite[][$1.7\\pm0.4\\times10^{-7}$ \\Mspyr; $3.7\\pm0.8\\times10^{-6}$ \\Mspyr]{km82,scuderi92}, analysing low-excitation \\ion{Fe}{2} lines in the ultraviolet \\cite[][$1-5\\times10^{-9}$ \\Mspyr]{hens82}, modeling the IR excess \\cite[][$6\\times10^{-7}$ \\Mspyr]{bc77}, and radio-flux limits \\cite[][$\\leq 2\\times10^{-7}$ \\Mspyr]{atw84}.  It is not unreasonable to suggest that the degree to which we understand A-type supergiants in general may be tested by Deneb.  Therefore, it is troubling that the published mass-loss estimates for Deneb vary so widely and this fact casts doubt on the what we can infer about the properties of more distant A-type supergiants. The proximity of Deneb allows one to more rigorously constrain its physical parameters and apply several different techniques to estimate its mass-loss rate.  For these reasons, in this paper we investiage Deneb's fundamental properties, and in particular its mass-loss rate.  To do this we employ a new stellar wind-model atmosphere package developed by \\cite{jpaphd} for use with the {\\tt PHOENIX}\\footnote{Also see http://phoenix.physast.uga.edu} general-purpose stellar and planetary atmosphere code \\cite[][and references therein]{hb99} and bring together new observations of Deneb at visible, infrared, millimeter, and radio wavelengths.  Our principal motivation for the constructing unified wind models has been that most analyses of early-type supergiants still suffer from severe limitations present in the model atmospheres which are employed.  Recent systematic analyses \\cite[]{vtg99} of the atmospheric parameters of A-type supergiants have employed LTE, line-blanketed, plane-parallel, static model atmospheres. These models do not include the effects (on the temperature structure and the synthetic spectrum) of metal line-blanketing and a spherically extended expanding atmosphere. Furthermore, most analyses do not allow for departures from LTE, particularly for species heavier than H and He, in computation of the opacity and line formation, which are especially important in low-density extended atmospheres.  \\cite{aph97} have developed models, employed by K99, which address many of these limitations including the solution of the spherical expanding transfer equation in the co-moving frame and the non-LTE treatment of hydrogen and helium.  Despite these important improvements, these models have had limited success in fitting spectral energy distributions (SEDs) due to the lack of metal line-blanketing in both the atmospheric structure and the synthetic spectrum.  Fundamental stellar parameters derived from line-profile fitting of hydrogen and helium lines alone may not be consistent with the SED.  Therefore, we believe that fitting both the line spectrum and the SED simultaneously should lead to more robust atmospheric parameters for A-type supergiants.    In \\S 2 we describe the spectrophotometric, spectroscopic, and interferometic observations used here. In \\S3 limb-darkening models and corrections to the angular diamter are discussed.  The effective temperature, reddening and other parameters constrained by the observations are discussed in \\S 4.  In \\S 5 we compare the spectrophotometic data from the UV to the radio with our synthetic spectral energy distributions.  The comparison of our models to portions of Deneb's line spectrum is discussed in \\S 6.  We summarize our results and conclusions in \\S 7.  Details on the model atmosphere construction and computation are left for Appendix A. ",
        "conclusions": "A stellar wind module has been developed for the \\phoenix\\ stellar atmosphere code for the purpose of computing atmospheric structures and detailed synthetic spectra of hot luminous stars.  These models incorporate a solution of the expanding spherically symmetric radiative transfer equation in the co-moving frame, a treatment of detailed metal line-blanketing in non-local thermodynamic equilibrium, and a convergence of the model temperature structure under the constraint of radiative equilibrium.  We have used the model atmospheres and synthetic spectra from this code to better illuminate what we do not understand about the spectral energy distribution, the mass-loss rate, and the fundamental stellar parameters of the A-type supergiant Deneb.  \\subsection{Millimeter and Centimeter Radiation Detections} We have reported the first detections of Deneb at millimeter and centimeter wavelengths with the SCUBA and the VLA.  These new data, in conjunction with IR photometery from 2 \\micron\\ to 60 \\micron, provide a better glimpse at the thermal continuum from Deneb's extended atmosphere.  The slope of the radio spectrum implies a stellar wind where hydrogen is partially ionized, a result reproduced by our models.  We find that the effects of line-blanketing on the synthetic millimeter and radio continua are significant and must be taken into when analyzing the continua of partially ionized winds.  \\subsection{Angular Diameter Measurement and Limb-darkening} Visibility measurements from the NPOI indicate a uniform-disk angular diameter for Deneb of $\\overline{\\theta_{\\rm UD}} = 2.40 \\pm 0.06\\ {\\rm mas}$.  Our models predict that the center-to-limb profile for Deneb in the $I$-band is significantly extended relative to hydrostatic model predictions.  We believe this limb profile causes the uniform disk angular diameter to be larger than the angular diameter $\\theta_{\\rm Ross} = 2.32 \\pm 0.06 {\\rm\\ mas}$ corresponding to $R_{\\rm Ross}$, the radius at the Rosseland mean optical depth $\\tau_{\\rm Ross}$ = 2/3.  The result is a correction to the uniform disk diameter which is less than unity, contrary to the limb-darkening correction made in the calculation of the effective temperature for more compact stars.  This limb profile prediction will hopefully be testable in the coming years as longer base-line optical/IR interferometers come on-line.  \\subsection{Interferometic Effective Temperature} Considering the uncertainites in the spectrophotometry, the interstellar extinction, and the corrected NPOI angular diameter, we find an inteferometric fundamental effective temperature of $T_{\\rm eff}^{\\rm Ross} = 8600 \\pm 500\\ {\\rm K}$ for Deneb. The biggest uncertainty in the value of the effective temperature is the uncertainty in the color excess, $E(B-V) = 0.09^{+0.04}_{-0.03}$. This effective temperature together with the distance estimates from {\\it Hipparcos} and Deneb's Cyg OB7 association membership, and theoretical evolutionary tracks yield: $R_{\\rm Ross} \\simeq$ 180 \\Rsun, $L \\simeq 1.6 \\pm 0.4 \\times 10^4$ \\Lsun, $M \\sim 20 -25$ \\Msun, and \\Logg\\ $\\simeq 1.3$.  \\subsection{Least-squares Spectrum Analysis} A least-squares comparison of 55 synthetic spectral energy distributions from 1220 \\AA\\ to 3.6 cm with the spectral energy distibution of Deneb provide estimates (1$\\sigma$) for the star's effective temperature and mass-loss rate of: $T_{\\rm eff}^{\\rm Ross} \\simeq 8420 \\pm 100$ K and $\\dot{M} \\simeq 8 \\pm 3 \\times 10^{-7}$ \\Mspyr.  The best fitting models to the SED indicate that the color excess toward Deneb is $E(B-V) = 0.06 \\pm 0.01$.  This effective temperature is consistent with the inteferometric fundamental effective temperature.  The smaller 1$\\sigma$ error in the least-squares value of $T_{\\rm eff}^{\\rm Ross}$ relative to interferometric value reflects the tighter contraint on the color excess from the least-squares fit.  However, both values of $T_{\\rm eff}^{\\rm Ross}$ are subject to systematic errors due to the poorly established interstellar extinction curve toward Deneb.  The least-squares analysis shows that, in addition to the effective temperature and mass-loss rate, the steepness of $\\beta$-law velocity field, the microturbulence, and the surface gravity all have a significant effect on the model temperature structure and ultraviolet spectral energy distribution.  It is interesting to point out that models with $\\beta$ outside the range $3.0\\pm1.0$ provide very poor fits to the UV spectrophotometry and that $\\beta$ values near 3.0 are not consistent with the radiation driven wind models of \\cite{alp97}, which predict $\\beta < 1$ for A-type supergiant winds.  \\subsection{Non-LTE Effects on the Line Spectrum} The mass-loss rate range derived from the least-squares analysis is reasonably consistent with that derived from high dispersion spectra of Deneb when non-LTE metal-line blanketing is treated, in particular for species \\myion{Fe}{i-iii}, \\myion{Mg}{ii}, and \\myion{Ca}{i-iii}. The primary effect of the non-LTE treatment is to enhance the ionization of these elements in the wind relative to LTE.  The non-LTE model ionization structure produces much lower column densities of \\ion{Fe}{2}, \\ion{Mg}{2}, and \\ion{Ca}{2} relative to LTE and causes strong lines to form deeper in the wind at lower velocities, in better agreement with the observations.  For the \\ion{Fe}{2} UV lines and the \\ion{Ca}{2} $H$ \\& $K$ lines, the non-LTE effects are striking.  The mass-loss rate inferred from the \\ion{Fe}{2} UV lines is 50 to 100 times larger in non-LTE than in LTE.  The synthetic $H$ \\& $K$ profiles show strong synthetic P-Cygni profiles in LTE which are not present in the data or in the non-LTE synthetic profiles.  \\subsection{Problems with \\ha\\ and Other Hydrogen Lines} We are unable achieve a reasonable fit to a typical \\ha\\ P-Cygni profile for Deneb with any model parameters over a reasonable range. The most prominent problem is that the model absorption component is always too strong.  While Deneb's \\ha\\ profile is variable, it is not hugely variable, and the model \\ha\\ profiles don't look like the observed profile at anytime. The wind model line profiles generally provide a better match to the higher Balmer series lines than hydrostatic model line profiles.  Both wind models and hydrostatic models match the higher Paschen series lines equally well.  Both types of models fail badly for the higher Pfund series lines.  These failures indicate that a spherically symmetric, expanding, steady state, line-blanketed, radiative equilibrium structure is not consistent with the conditions under which a typical \\ha\\ and the higher Pfund lines form.  The strong lines of \\ion{Fe}{2}, \\ion{Mg}{2}, and \\ion{Ca}{2} that we can reproduce with our models originate from the ground state or low-exitation levels and apparently are not sensitive to the conditions which affect the 2$\\rightarrow$3 transition of hydrogen. Our inability to fit a typical \\ha\\ profile from Deneb is very bothersome because \\ha\\ is the most commonly used mass-loss rate diagnositic for A-type supergiants.  \\subsection{Mass-loss Rate} The thermal millimeter and radio continua of Deneb when compared with our models provide both lower and upper limits to its mass-loss rate. The detection of radiation at 3.6 cm, far in excess of what is expected for no mass-loss, provides the lower limit at about $10^{-7}$ \\Mspyr. The $3\\sigma$ upper limit to the 870 \\micron\\ flux provides the upper limit at about $10^{-6}$ \\Mspyr.  If Deneb's extended atmosphere departs significantly from spherical symmetry, this would systematically bias our estimates for the mass-loss rate.  However, there is no evidence for deviations from spherical symmetry from instrinsic polarization studies, either in the continuum or in the spectral lines \\hb, \\hg, \\ion{Ca}{2} $H$ \\cite[][and references therein]{notpolarized}.  The line spectrum of Deneb provides mainly upper limits to the mass-loss rate.  While the observed P-Cygni character of the \\ha\\ and \\ion{Mg}{2} $h$ and $k$ lines are clear indicators of the stellar outflow, a hydrostatic model with no mass loss is really not that bad a fit to most of the UV line spectrum.  It is the more subtle characteristics of the UV lines, their widths and desaturated profiles which show evidence of the velocity field.  Many weak metal lines in the optical spectrum are also very well matched by models with no mass-loss and all our models with \\mdot $> 10^{-8}$ \\Mspyr\\ predict profiles for these lines which on average tend to be washed out and too shallow.  Certainly the column densities in many of these transitions are too high.  While non-LTE ionization effects solve this problem for some lines, others are still problematic.  A line-by-line study with a complete grid of models which include \\myion{Ti}{i-iii}, \\myion{Cr}{i-iii}, and other metals in non-LTE is needed to analyze this issue fully.  We have model atoms for these species and will investigate this in future work.   While previously published estimates for Deneb's mass-loss rate range over more than three orders of magnitude, we can with confidence reduce this uncertainty to one order of magnitude, but not much better.  Reducing this uncertainty further requires a better match simultaneously to the spectral energy distribution and the spectral lines. This may require removing several of our simplifying assumptions such as radial symmetry, homogeniety (filling-factor equal to unity), time-independence, and no mechanical dissipation.  We suggest that before the Wind Momentum-Luminosity Relationship, thus far based solely on Balmer line fits, can be firmly established for A-type supergiants, researchers should work to toward firmly establishing fundamental parameters and stellar wind properties of Galactic A-type supergiants by checking that models are consistant with both the spectral energy distributions and the line spectra."
    },
    "0201/hep-th0201180_arXiv.txt": {
        "abstract": "We discuss the problem of the motion of classical strings in some black hole and cosmological spacetimes. In particular, the null string limit (zero tension) of tensile strings is considered. We present some new exact string solutions in Reissner-Nordstr\\\"om black hole background as well as in the Einstein Static Universe and in the Einstein-Schwarzschild (a black hole in the Einstein Static Universe) spacetime. These solutions can give some insight into a general nature of propagation of strings (cosmic and fundamental) in curved backgrounds. ",
        "introduction": "\\label{intro}  Fundamental string theory is undoubtedly the most serious candidate for unification of gauge interactions with gravity \\cite{green}. Its effects should clearly be visible in extremely high gravitational fields of black holes and in the early universe. It is not an easy task to study quantum string propagation in these background fields and this gives motivation to study the motion of classical strings in these fields first in order to catch uo some really ``stringy'' properties of a quantum theory. On the other hand, classical motion of strings gives an appropriate formalism to study the dynamics of cosmic strings which appear naturally in GUT models \\cite{vilenkinbook}. This is why we will study the classical motion of strings in some black hole and cosmological spacetimes.  Classical motion of strings which evolve in curved spacetimes can be described by a system of the second-order non-linear coupled partial differential equations \\cite{vesan,integr}. The non-linearity of these equations gives a complication which leads to their non-integrability and possibly chaos \\cite{ott}. It is well-known that various types of nonlinearities appear in Newtonian as well as relativistic systems and so they can deliver chaos. On the other hand, some types of non-linear equations can be integrable and their solutions are not chaotic. It seems that theory of relativity is ideal to produce chaotic behaviour since their basic equations are highly non-linear. However, the problem is not as easy as one could think of, because most of the systems under study possess some symmetries which simplify the problem. This also refers to a single particle obeying either Newtonian or relativistic equations. Simply, a single particle which moves in the gravitational field of a source of gravity cannot move chaotically. However, two particles which form a 3-body system including the source can move in a chaotic way, though still not for all possible configurations.  Admission of extended objects such as strings gives another complication which, roughly, can be compared to the fact that now we have a many-body system which can obviously be chaotic on the classical level. An extended character of a string is reflected by the equations of motion which become a very complicated non-linear system from the very beginning. Thus, no wonder chaos can appear for classical evolution of strings around the simplest sources of gravity such as Schwarzschild black holes. This, in fact, was explicitly proven \\cite{lar94,frolar}. However, in a similar way as for other types of non-linear sets of equations, there exist integrable configurations. The investigation of such explicit configurations can give an interesting insight into the problem of the general evolution of extended objects in various sources of gravity. Of course, it is justified, provided we do not consider back-reaction of these extended objects onto the source field, i.e., if we consider test strings in analogy to test particles which do not ``disturb'' sources' gravitational fields.  Studies of exact configurations can give big insight into the problem. One useful example is when unstable periodic orbits (UPO) appear. Their emergence becomes a signal for a possible chaotic behaviour of the general system \\cite{frankel}.  The task of this paper is to study some exact configurations for strings moving in simple spacetimes of general relativity. Unfortunately, for strings, the main complication refers to their self-interaction reflected in the equations of motion by a non-zero value of tension (tensile strings). However, one is able to study simpler extended configurations for which tension vanishes called null (tensionless) strings \\cite{null,venico,LS96}. Their equations of motion are null geodesic equations of general relativity appended by an additional `stringy' constraint. Many exact null string configurations in various curved spacetimes have already been studied \\cite{LS96,kar,dablar,mar2,alex,porfyriadis98,kuirokidis99,jassal}. One of the advantages of the null string approach is the fact that one may consider null strings as null approximation in various perturbative schemes for tensile strings \\cite{venico,alex1,alex2,alex3}.  In section II we present tensile and tensionless string equations of motion. In Section III we obtain exact null string configurations both in Reissner-Nordstr\\\"om and Schwarzschild spacetime while in Section IV we derive string configurations in static Einstein Universe. In Section V we discuss the evolution of strings in Einstein-Schwarzschild (Vadiya) Universe. In Section VI we discuss our solutions. ",
        "conclusions": "\\label{concl}  In this paper we have found some exact string configurations in black hole and cosmological spacetimes which apply both for fundamental and for cosmic strings. We generalized previously found solutions of Ref. \\cite{dablar} for a \"cone\" string and for a string moving on the photon sphere into a Reissner-Nordstr\\\"om spacetime which is also related to the discussion of the behaviour of strings in this spacetime given in refs. \\cite{LS93,LS96,LarS96}. We also generalized an event horizon solution and presented a Cauchy horizon solution for the Reissner-Nordstr\\\"om spacetime. We found a solution for a null string moving around the Einstein Static Universe and two completely new solutions for strings evolving in the Einstein-Schwarzschild spacetime (a black hole in the Einstein Static Universe).  Firstly, we briefly presented formalism which allowed to take the limit of null strings in an appropriate action. Then, we studied the evolution of strings in Reissner-Nordstr\\\"om, Einstein Static and Einstein-Schwarzschild spacetimes. The exact configurations we found can be grouped geometrically into a couple of classes. There is a class of solutions which describe null strings residing on the null surfaces of these spacetimes, i.e., event and Cauchy horizons. There is also a class of solutions which describe strings sweeping out the light cones of a particular spacetime. Another class is for strings which reside on the surface of the photon sphere (an unstable periodic orbit for zero point particles). This class exists both in Schwarzschild/Reissner-Nordstr\\\"om spacetimes and, in an adapted form, in the Einstein Static Universe, but not in the Einstein-Schwarzschild spacetime.  As far as the physical properties are concerned we found that some of our solutions are unstable (for instance, a string on the photon sphere in Reissner-Nordstr\\\"om spacetime) and some are stable (e.g., a string on the event horizon). According to Ref. \\cite{multi}, multistring solutions appear whenever the world-sheet time $\\tau$ is a multi-valued function of the physical time and they are possible, for instance, in the positive cosmological constant models such as the de Sitter space. In our paper only the Einstein Static Umiverse admits a positive cosmological constant and because of that one should perhaps expect some multistring solutions admissible. However, our solutions of Sections IV and V do not possess this property. On the other hand, some of our solutions (a string on the photon sphere (Eqs. (\\ref{jeden})-\\ref{troje})) and a string in the Einstein Static Universe (Eqs. (\\ref{SPI1})-(\\ref{SPI4})) have an invariant string size described by multiply covering azimuthal angle because of an infinite domain of the timelike string coordinate $\\tau$.  The existence of the photon sphere, i.e., an unstable periodic orbit (UPO), together with other special solutions suggests that a general evolution of a tensile (or perhaps even a null) string in these simple curved backgrounds is chaotic. This statement is obviously true for Schwarzschild spacetime \\cite{frolar}, and the solutions we have found are straightforward generalizations of exact configurations in Schwarzschild spacetime.  The results we gained can give some insight into the nature of motion of strings in extremely high gravitational fields of black holes and in the early universe in fully quantum string theory."
    },
    "0201/astro-ph0201054_arXiv.txt": {
        "abstract": "We provide a set of microlensing event rate maps for M31, the Andromeda Galaxy. Rates for M31 microlensing were calculated on the basis of a four component model of the lens and source populations: disk and bulge sources lensed by bulge, disk, M31 halo and Galactic halo lenses.  We confirm the high rate gradient along the minor axis of M31 due to a dark halo of lenses. Furthermore, we compute the timescale distributions of events, for both Einstein times and full-width at half-maximum times.  We explore how the rate contours and the timescale distributions can be used to measure the shape and extent of the microlensing halo.  With one year of twice--weekly sampling, or three observing seasons, a halo MACHO fraction as small as 5\\% can be detected with modest ground based telescopes. ",
        "introduction": "Gravitational microlensing as a means to detect compact objects in the Galactic halo was first considered by Paczy\\'nski (1986), but the basic idea is much older (Einstein 1936). This suggestion was realized as results from two surveys of microlensing events towards the Large and Small Magellanic Clouds (LMC and SMC) (Alcock et al.~2000; Ansari et al.~1996).  These were consistent with a significant, but subdominant, contribution of microlensing masses to the Galactic dark matter halo. Nonetheless, these conclusions are still controversial, and the identity and location of the microlensing masses are still mysterious.  A decade ago, M31 was suggested as a promising venue where galactic microlensing might be explored in ways advantageous and distinctive from that in and around the Galaxy (Crotts 1992).  Several papers (Jetzer 1994; Han \\& Gould 1996; Baltz \\& Silk 2000; Kerins et al.\\ 2001) have confirmed that a substantial microlensing signal can be expected.  Two collaborations, MEGA (preceded by the VATT/Columbia survey) and AGAPE have produced a number of microlensing event candidates involving stars in M31 (Crotts \\& Tomaney 1996, Ansari et al.~1999, Auri\\`ere et al.~2001, Uglesich 2001, Calchi Novati et al.~2002).  Here we show the potential for these and future surveys to settle some of the outstanding questions regarding microlensing in spiral galaxies.  This is the second paper in a series.  Paper I (Gyuk \\& Crotts 2000) provides optical depth maps for M31. While these are useful tools for certain purposes they are unfortunately not directly measurable. Event rates on the other hand {\\em are} directly measurable and hence their magnitude and variation across the face of M31 are more meaningful in planning and evaluating surveys of microlensing. In this paper we extend Paper I to include event rate maps, both total and also differential rates with respect to the event timescale.  This paper is organized as follows.  In \\S\\ref{model} we briefly discuss the M31 models we used, including disk, bulge and halo components.  Following this we present rate maps for various halo models, including the self lensing contribution, in \\S\\ref{ratemaps}.  In \\S\\ref{timescales} we provide differential rate distributions as a function of two different timescale measures, and we discuss how cuts in timescale can be used to separate self lensing from a MACHO halo contribution to the lensing rate.  We conclude with a discussion of measuring lens masses and halo properties in \\S\\ref{discussion}. ",
        "conclusions": "\\label{discussion}  We have provided maps of the microlensing event rate towards M31 in a model including M31 disk and bulge sources, and lenses from the M31 disk and bulge, as well as lenses in the dark halos of M31 and the Milky Way.  We have varied the parameters of the dark halo, namely the core radius and the flattening, and studied the effects on the rate contours.  The core radius affects the rates most in the galaxy center but these changes are partially obscured by the high rate of self lensing events there.  The flattening does have a significant effect, especially along the major axis.  The normalization of the rate contours away from the bulge gives a measure of the halo MACHO fraction.  The easily measured timescale parameter is the full width at half maximum, but the more physical timescale is the Einstein time.  Knowing the Einstein times allows a much more accurate measurement of the lens mass, thus it is worth the extra trouble to try to measure the Einstein times of the detected events.  We have shown that a microlensing halo in M31 should be clearly distinguishable from self lensing if an appreciable event rate away from the M31 bulge is measured.  We have explored the use of cuts in event timescale to separate the self lensing component from the halo lensing component of the event rate, and found that this helps near the bulge, but not further away.  We have quantified the level of halo that is detectable, and found that a marginal detection of a 5\\% microlensing halo would be possible in three seasons of ground--based observations.  Higher halo fractions can be detected more convincingly of course."
    },
    "0201/astro-ph0201262_arXiv.txt": {
        "abstract": "GRB data accumulated over the years have shown that the distribution of their time duration is bimodal.  While there is some evidence that long bursts are associated with star-forming regions, nothing is known regarding the class of short bursts. Their very short timescales are hard to explain with the collapse of a massive star, but would be naturally produced by the merger of two compact objects, such as two neutron stars (NS), or a neutron star and a black hole (BH).  As for the case of long bursts, afterglow obervations for short bursts should help reveal their origin. By using updated population synthesis code calculations, we simulate a cosmological population of merging NS-NS and NS-BH, and compute the distribution of their galactic off-sets, the density distribution of their environment, and, if indeed associated with GRBs, their expected afterglow characteristics. ",
        "introduction": "Since data on GRBs started to accumulate over the past two decades, it was recognized that their time distribution appears to be bimodal\\footnote{Note however the recent claims about the possibility that a third class may exist (Balastegui et al. 2001).}, with about 25\\% of bursts having a short duration, of mean $\\sim 0.2$ sec, and the rest having a much longer duration, of mean $\\sim 20$ s (Mazets et al. 1981; Hurley et al. 1992; Kouveliotou et al. 1993; Norris et al. 2000).  The separation between the two classes appears to be around 2 s.  While the two classes of bursts seem to have similar (isotropic) spatial distributions, they differ in several other respects. Short bursts tend to have harder spectra than long bursts (Kouveliotou et al. 1993; Dezelay et al. 1996), and about 20 times less fluence\\footnote{Lee \\& Petrosian (1997) showed indeed that there is a highly significant positive correlation between the burst fluence and duration.} and it has also been presented some evidence that their number-intensity distribution (Belli 1997; Tavani 1998) differs from that of the longer class.  A recent analysis by Norris et al. (2000) of the temporal properties of the bursts, such as the distribution of number of pulses per burst, pulse width, and intervals between pulses, clearly slowed that the two classes of long and short bursts are disjoint.  The existence of two distinct populations of GRBs might very well be an indication of the presence of two distinct types of progenitors. The currently favoured GRB models can be divided into two classes: models involving mergers of two compact objects (Goodman 1986; Eichler et al. 1989; Paczynski 1991; Narayan et al. 1992; Meszaros \\& Rees 1992; Katz \\& Canel 1992), and models involving the collapse of a massive star (Woosley 1993; Paczynski 1998; MacFadyen \\& Woosley 1999; Vietri \\& Stella 1998).  According to the internal shock model (see e.g. Piran 1999 for a review) for the production of the observed $\\gamma$-ray emission, the duration of the event, whatever it is, is very likely to be a direct measure of the time interval during which the powering engine is active.  Simulations of mergers of two compact objects have shown that the duration of the neutrino-driven wind possibly producing the GRB (Ruffert \\& Janka 1999) is less than a second. On the other hand, the relativisitc outflow generated by the collapse of a massive star (MacFadyen \\& Woosley 1999) can last several tens of seconds.  Therefore, if one were to associate the two classes of GRBs with two classes of models, it would be natural to associate long GRBs with the collapse of massive stars and short ones with mergers of two compact objects.  Traditionally, it is considered that an important way to test the above assumption and possibly distinguish between the two classes of models is by determining the location in which the bursts occur. Massive stars have short lifetimes, and therefore are expected to die close to where they are born, that is in dense and dusty environments. On the other hand, compact objects receive kicks when they are born, and are therefore expected to travel far from their birthplaces.  Along these lines, several methods have been proposed to constrain the GRB location and the characteristics of their environment.  The long-lived remnants resulting from the interaction between GRBs and their afterglows with the surrounding medium can be identified in nearby galaxies based on their spectral signatures (Wang 1999; Perna, Raymond \\& Loeb 2000; Perna \\& Raymond 2000) or their dynamical interaction with the medium (Efremov, Elmegreen \\& Hodge 1998; Loeb \\& Perna 1998, Ayal \\& Piran 2001), therefore allowing a close study of the GRB sites.  On the much shorter time scale during which the afterglow propagates in the medium, several independent analyses can be made.  When multifrequency data are available, then a determination of the various break frequencies and the peak flux in the afterglow spectrum (Sari, Piran \\& Narayan 1998) allows to constrain the burst parameters; this has been done in several cases by a number of authors (e.g. Wijers \\& Galama 1999; Panaitescu \\& Kumar 2001a).  Time variability of absorption lines due to the gradual photoionization of the medium by the X-ray UV radiation is also sensitive to the type of environment (Perna \\& Loeb 1998; B\\\"ottcher et al. 1998; Lazzati et al. 2001), as it is the time delay between the $\\gamma$-ray emission and the onset of the afterglow (Vietri 2000).  So far, afterglow observations have only been possible for long GRBs\\footnote{See however Lazzati, Ramirez-Ruiz and Ghisellini 2001 for a possible detection in the BATSE data.}. For this class of bursts, there has been mounting evidence that they are associated with the collapse of massive stars. Bloom, Kulkarni \\& Djorgovski (2001) compared the observed offset distribution of 20 GRBs with the theoretical predictions of two models, one which is representative of collapsars and promptly bursting binaries (such as binaries in which the black hole merges with the helium core of an evolved star during a common envelope phase), and another representative of delayed merging remnants.  They found that the latter population can be ruled out to a high confidence level. Their conclusions are strengthened by the observed correlation of GRB locations with the UV light of their hosts, which is strongly suggestive of the occurrence of GRBs in star-forming regions (see also Sahu et al. 1997, Kulkarni et al. 1998, Fruchter et al. 1999, Kulkarni et al. 1999).  An analysis of the evolution of the X-ray prompt emission of GRB 980329 and GRB 780506 by Lazzati \\& Perna (2001) has led to similar conclusions. An intriguing hint towards the connection of GRBs with the collapse of massive stars has been provided by the presence of a bump, interpreted as an underlying supernova component, in the light curve of GRB 900326 (Bloom et al. 1999) and GRB 970228 (Reichart 1999, Galama et al. 2000). The evidence for extinction by dust of some burst afterglows has provided further support to the GRB-collapsar connection. Finally, the recent detection of an iron line in the afterglow spectrum of 5 GRBs (Piro et al. 1999, 2000; Yoshida et al. 1999; Antonelli et al. 2000; Amati et al. 2000) provides evidence for the presence of dense matter in the vicinity of the burst sites (Vietri et al. 1999; Weth et al. 2000; Lazzati, Covino \\& Ghisellini 2000a).  The association between GRBs and collapsars, once well established, would have very important implications for our understanding of the star formation history in the universe (Blain \\& Natarajan 1999).  This huge wealth of information on the class of long GRBs has been gathered as a result of the afterglow observations.  No similar types of studies have been possible for the short bursts so far.  At this point, optical and radio searches have been performed only for the 4 bursts that were well pin pointed by the Interplanetary Network (Hurley et al. 2001), but the search did not lead to detections.  The situation is however going to change with HETE II and then {\\em Swift}, which will provide quick arc-minute localizations; longer-wavelength follow-ups will then be possible also for this class of short events, therefore allowing to do the same type of science for them as well.  As already mentioned, a strong candidate for the class of short bursts is provided by the coalescence of two compact objects, whose timescales and energetics are compatible with those inferred for the class of short GRBs.  Therefore an analysis of the afterglow properties that such a population would have is needed. This is the goal of this work. More specifically, in this paper we use updated population synthesis code calculations to study the afterglows that a cosmological population of GRBs due to mergers of compact objects should have. We consider two classes of progenitors: double neutron stars, and neutron star -- black hole. The former population includes the new class of short-lived neutron stars identified by Belczynski \\& Kalogera (2001), Belczynski, Kalogera \\& Bulik (2001c) and Belczynski, Bulik \\& Kalogera, (2001a). As it will be discussed in the following, this population has a very short lifetime and it dominates the merger rates. The environment of this population would therefore be more similar to that of collapsars and helium star-black hole mergers. On the other hand, the population of NS-BH binaries has a much longer lifetime, and therefore it has time to travel further away from its birthsite. More generally, we want to point out that, whereas the motivation of this paper has been a detailed study of merger events in relation to the class of short bursts, however our results regarding the new class of tight NS-NS binaries can be potentially relevant also for the class of long GRBs. In fact, as discussed above, some of the evidence on the connection of long GRBs with massive stars is based on the association of long GRBs with star forming regions, but this would also be the case for the new population of tight NS-NS binaries.  To generate our simulated set of data, we incorporate the results of the population synthesis code {\\em StarTrack} by Belczynski et al.\\ (2001c) within the context of a cosmological model which, with the help of a Monte Carlo type approach, accounts for (i) the redshift distribution of the merger events, (ii) the mass distribution of the galaxies where the events occur (which is a function of redshift) using a Press-Schechter type formalism (Press-Schechter 1974); (iii) the redshift dependence of the probability that a certain merger occurs at a given position within a galaxy of a given mass. This last effect is particularly important for the population of NS-BH binaries, whose lifetime can be comparable with the Hubble time, and it has not been considered so far. A byproduct of our computation is the distribution of the offsets that merging binaries have from the centers of their host galaxies.  For each merger event, the density of the surrounding medium is also determined within the simulation itself.  The other parameters that are needed to compute the afterglow intensity for each event are randomly drawn from distributions which have the typical values found in the afterglow modelling of long GRBs.  The paper is organized as follows: in \\S 2 we describe the various elements of the computation, which include the population synthesis code, the new population of double NS-NS, the galaxy potential and its density profile, and the afterglow modelling. The results of the simulation of the data are presented in \\S 3, while \\S 4 is devoted to a discussion with conclusions. ",
        "conclusions": "We have studied the properties that a population of GRB events due to mergers of compact objects should have with special emphasis on the related afterglows.  By using a Monte Carlo type of approach, our simulations take into account the mass distribution of the host galaxies as a function of redshift, as well as the redshift evolution of the probability distributions for the location of the mergers within galaxies of various mass (cfr. Fig. 1abc). This last effect needs to be taken into account especially for binaries whose merger time can be comparable with the Hubble time for a sizeable fraction of them (such as the NS-BH population).  Our population of double neutron star binaries includes the new groups of short-lived binaries identified by Belczynski \\& Kalogera (2001) and Belczynski et al.\\ (2001a, 2001c), which dominate the merger rates. Therefore our results regarding this population differ from previous studies on the same subject, and the derived distributions trace rather closely the star forming regions in the disk. The densities in which they occur are typical ISM densities; hence their afterglows, even though dimmer due to the smaller energy released, should still be observable with current X-ray instruments for a large fraction of them. Afterglows produced as a result of NS-BH mergers are even dimmer, but most of them should still be detectable if observed within the first few hours.  Whereas the NS-NS class of candidate GRB progenitors might not be distinguishible from that of collapsars and of other promptly-bursting binaries simply on the basis of their location within the host (and the consequent intensity of their afterglows), there are however other signatures, such as the presence of an underlying supernova explosion, or of iron lines in the afterglow spectrum (see Introduction) that, while naturally associated with a collapsar, would be hard to explain within the NS-NS merger scenario\\footnote{It should however be noted that ``SN bumps'' can also be explained as the result of dust echos (Esin \\& Blandford 2000), and that models that explain iron lines (e.g. Vietri \\& Stella 1998) require a SN that took place a few months before the GRB, which is inconsistent with the presence of the SN bumps.}.  On the other hand, the observational properties of the NS-BH binary population that we have studied here, and which depend on the location within the hosts (such as offsets, densities, and density-dependent afterglow quantities), differ significantly from those of the NS-NS population (which, as far as the location is concerned, could well be representative also of collapsars). Even if it is not possible to infer all the parameters $E\\;,n,\\;\\xi_e,\\;\\xi_B$ at once (which requires that $\\nu_c$, $\\nu_m$, $\\nu_a$ and $F_{\\rm peak}$ be all measured), a {\\em comparison} of the distributions for, say, $F_{\\rm peak}$ or $t_c$ for the class of long and that of short bursts would provide strong constraints on whether their progenitors actually belong to two different classes of progenitors, one which is short-lived and the other which is long-lived.  The population synthesis code that we used in all the simulations is the {\\em StarTrack} code by Belczynski et al. (2001c), operated in its ``standard'' mode, where the best values of all the parameters are chosen.  A complete parametric study of how the distributions for the location within a galaxy of a given mass change when all the model parameters are varied to their extremes is being performed elsewhere (Belczynski et al.\\ 2001b). The main results are that the merger site distribution has the strongest dependence on the prescriptions for the mass transfer and the common envelope efficiency, and it is also rather dependent on the maximum allowed NS mass and the kick velocity. It is not very dependent on the assumed cosmology. Regarding the new class of short-lived binaries identified by Belczynski \\& Kalogera (2001) and Belczynski et al. (2001a, 2001c), it is found that 81\\% of them contribute to the NS-NS population in our standard model described in \\S 2.1. The parametric study shows that  the highest contribution (98\\%) from this population is obtained for small kicks, while the smallest (28\\%) for very low CE efficiency. We stress once again that our results on GRBs from NS-NS mergers strongly rely on the presence of this short-lived population, whose presence is based on the assumption that low-mass helium stars can survive the CE phase. This will have to be tested through detailed hydrodynamical simulations.  In this work, our main interest  has been to incorporate the results of the population synthesis code within a proper cosmological context, and study the expected afterglows if GRBs are indeed associated with mergers of two compact objects. This is particularly relevant for the population of short bursts, whose very short timescales are hard to account for with the collapse of a massive star, while being naturally associated with mergers of two compact objects."
    },
    "0201/astro-ph0201438_arXiv.txt": {
        "abstract": "A maximum-likelihood method is presented for estimating the power spectrum of anisotropies in the cosmic microwave background (CMB) from interferometer observations.  The method calculates flat band-power estimates in separate bins in $\\ell$-space, together with confidence intervals on the power in each bin. For multifrequency data, the power spectrum of the other foreground components may also be recovered.  Advantage is taken of several characteristic properties of interferometer data, which together allow the fast calculation of the likelihood function. The method may be applied to single-field or mosaiced observations, and proper account can be taken of non-coplanar baselines. The method is illustrated by application to simulated data from the Very Small Array. ",
        "introduction": "\\label{intro}  Interferometers have proven themselves to be valuable tools in observing anisotropies in the cosmic microwave background (CMB). They are inherently insensitive to anisotropies in atmospheric emission on scales larger than the beam and to ground spillover, and the fact that they sample Fourier space directly make interferometers ideal instruments for measuring the CMB power spectrum.  A method for calculating the maximum-likelihood CMB power spectrum directly from the complex visibility data produced by an interferometer was first discussed by Hobson, Lasenby \\& Jones (1995) (hereafter HLJ95). The method was illustrated by assuming the CMB anisotropies to be described by a Gaussian autocorrelation function, but the technique is easily modified so that the CMB power spectrum is parameterised in terms of flat band-powers in separate bins in $\\ell$-space.  Indeed, the modified HLJ95 algorithm was used to calculate flat band-power estimates of the CMB power spectrum in two spectral bins at $\\ell \\approx 410$ and $\\ell \\approx 590$ (with bin-width $\\Delta\\ell \\approx 90$) for the Cosmic Anisotropy Telescope (CAT), which is a 3-element interferometer (Scott et al. 1996; Baker et al. 1999)  Following the early success of the CAT, a new generation of CMB interferometers have been built, and have recently made high-sensitivity observations of the CMB.  These experiments include the Very Small Array (VSA) (Jones 1997; Jones \\& Scott 1998), the Degree Angular Scale Interferometer (DASI) (Leitch et al. 2001; Halverson et al. 2001) and the Cosmic Background Imager (CBI) (Pearson et al. 2000; Padin et al. 2001). Although the detailed design of these experiments is different in each case, the basic principles underlying their operation are the same.  In particular, these instruments have a larger number of antennas than the CAT (for example, the VSA has 14 horns) and more sensitive detectors. These specifications enable the accurate measurement of the CMB power spectrum over a wide range of angular scales.  For example, the VSA in its `compact' configuration (see section~\\ref{application}) measures the CMB power spectrum in 10 independent bins of width $\\Delta\\ell \\approx 90$ from $\\ell \\approx 80 - 950$.  Some discussion of how to obtain flat band-power estimates of the CMB power spectrum from this new generation of interferometer experiments has been presented by White et al. (1999a) and White et al. (1999b), mainly in connection with the analysis of DASI data. Indeed, the techniques outlined by these authors have been applied to the analysis of DASI and CBI observations.  In particular, the DASI experiment has recently produced an accurate determination of the CMB power spectrum in 9 spectral bins of width $\\Delta\\ell \\approx 80$ in the range $\\ell \\approx 100 - 900$ (Halverson et al. 2001), whereas the CBI has measured flat band-power estimates of the CMB power spectrum in two spectral bins at $\\ell \\approx 600$ and $\\ell \\approx 1200$, with a spectral resolution of $\\Delta\\ell \\approx 400$ (Padin et al. 2001).  The increase in both the amount of visibility data and the number of independent spectral bins means that the computational burden of performing a likelihood analysis of the new generation of interferometer experiments is considerable. It is therefore of interest to investigate fast methods of calculating the likelihood function for interferometer observations of the CMB.  In this paper, we extend the discussion given by HJL95 and White et al. (1999a, 1999b) and present a complete description of a maximum-likelihood technique for calculating flat band-power estimates of the CMB power spectrum from interferometer data.  Following HJL95, the technique allows for contributions to the visibilities from Galactic foreground emission, which can be separated out if multifrequency data are available. In terms of the computational algorithm, we also present some straightforward devices for speeding up the calculation of the likelihood function, which take advantage of certain useful properties of interferometer data.  Given the facility of fast evaluation of the likelihood function, we then investigate the relative merits of obtaining the maximum-likelihood flat band-power estimates, and their corresponding confidence intervals, either by direct evaluation the likelihood distribution, or by using traditional numerical maximisation techniques.  Our formalism makes no special assumptions specific to a particular experiment, and so can be applied to data from any CMB interferometer, including observations of mosaiced fields, and proper account is taken of non-coplanar baselines. The method is illustrated by applying it to simulated data from the Very Small Array (VSA). The application of the technique to real VSA observations will be presented in a forthcoming paper. We note, in passing, that a maximum-entropy map-making method for interferometer observations, which can simultaneously deconvolve the interferometer beam and separate CMB and Galactic foreground emission, is discussed by Maisinger, Hobson \\& Lasenby (1997). ",
        "conclusions": "\\label{conc}  In this paper, we have investigated maximum-likelihood methods for estimating the power spectrum of the cosmic microwave background (CMB) from interferometer observations. In particular, we have considered the computational efficiency of several techniques for obtaining flat-band power estimates in a number of spectral bins, together with confidence limits on the power in each bin. For multifrequency data, one may also estimate the flat band-powers of Galactic foreground emission in each spectral bin. The methods developed may be applied to single-field or mosaiced observations, and take proper account of non-coplanar baselines.  The sparse nature of the covariance matrix for visibility data allows the use of sparse matrix techniques which significantly reduce the computational burden of evaluating the likelihood function, as compared with standard dense matrix routines. This enables the maximum-likelihood solution to be obtained more efficiently by numerically maximising the likelihood using only function values, as opposed to more traditional techniques, such as the Newton--Raphson algorithm, which rely on gradient and curvature information. We also find the covariance matrix of the errors on the parameters is most efficiently calculated using a numerical second-differencing approach, rather than direct calculation of the analytic expression for the curvature matrix.  The speed with which the likelihood function is calculated may be further increased by making use of the fact that only a small fraction of the total number of observed visibilities are sensitive to the flat band-power in any one spectral bin. Using this reduced data-set for each parameter enables very fast evaluation of the likelihood function along the `axes' in parameter space. Indeed, we find that performing independent line-maximisations of the likelihood function along these parameter directions, and iterating until convergence, provides the computationally most efficient way of obtaining the maximum-likelihood solution.  If the spectral bins are chosen to be sufficiently wide that the flat band-powers are quasi-uncorrelated parameters, one may dispense with the Gaussian approximation to the likelihood function near its peak, and obtain (generally asymmetric) confidence intervals on each parameter by calculating the likelihood function along each  parameter direction through the maximum likelihood point. This is best achieved by performing a single-to-noise eigenmode rotation for each parameter. Indeed, this method can itself be used to arrive at the maximum-likelihood solution. This is illustrated by application to simulated observations by the Very Small Array in its `compact' configuration. If multifrequency data are available, this technique may also be used to marginalise over the flat band-power of foreground Galactic emission in each spectral bin.   \\subsection*"
    },
    "0201/astro-ph0201112_arXiv.txt": {
        "abstract": "During the early stages of galaxy evolution, the metallicity is generally low and nearby metal-poor star-forming galaxies may provide templates for primordial star formation. In particular, the dust content of such objects is of great importance since early molecular formation can take place on grains. To gain insight into primeval galaxies at high redshift, we examine the dust content of the nearby extremely low-metallicity galaxy SBS 0335-052 which hosts a very young starburst ($\\la 10^7$ yr). In young galaxies, the dust formation rate in Type II supernovae governs the amount of dust, and by incorporating recent results on dust production in Type II supernovae we model the evolution of dust content. If the star-forming region is compact ($\\la 100$ pc), as suggested by observations of SBS 0335-052, our models consistently explain the quantity of dust, far-infrared luminosity, and dust temperature in this low-metallicity object. We also discuss the H$_2$ abundance. The compactness of the region is important to H$_2$ formation, because the optical depth of dust for UV photons becomes large and H$_2$ dissociation is suppressed. We finally focus on implications for damped Ly$\\alpha$ systems. ",
        "introduction": "How dust forms and evolves in primordial galaxies needs to be considered before we can understand the chemical and thermodynamical evolution of a metal-poor interstellar medium (ISM). Dust grains absorb stellar light and reemit it in far infrared (FIR)\\footnote{For SBS 0335-052, large energy is emitted in mid-infrared because the dust temperature is high (Dale et al.\\ 2001). But in this Letter, we use the term ``FIR emission'' for the thermal emission from dust grains.}. Indeed, the FIR spectral range represents a unique opportunity to study the dust properties and distribution, and dedicated space missions are planned in these wavelength bands (ASTRO-F\\footnote{http://www.ir.isas.ac.jp/ASTRO-F/index-e.html}, SIRTF\\footnote{http://sirtf.caltech.edu/}, Herschel\\footnote{http://astro.estec.esa.nl/First/}, etc.). Dust grains also drastically accelerate the formation rate of molecular hydrogen (H$_2$), expected to be the most abundant molecule in the ISM. Hydrogen molecules emit vibrational-rotational lines, thus cooling the gas. The cooling rate of a galaxy consequently depends strongly on the dust content and its effect on H$_2$ abundance.  The production of dust is expected at the final stages of stellar evolution, and Type II supernovae (SNe II) are the dominant source for the production of dust grains in young star-forming galaxies. The formation in stellar winds from evolved low-mass stars can also contribute considerably, but the cosmic time is not long enough for such stars to evolve at high redshift ($z$). However, dust is also destroyed by SN shocks (Jones, Tielens, \\& Hollenbach 1996). The detailed modelling of dust evolution in star-forming regions therefore requires an accurate treatment of both types of processes.  Here we have developed a model for the evolution of dust content in primordial galaxies. To compare our model with observations, we need age information and the FIR energy distribution. For high-$z$ primordial galaxies, present experiments cannot determine these quantities accurately enough. Nevertheless, because of their low metallicity and active star formation, blue compact dwarf galaxies (BCDs) may be viable candidates for nearby primeval systems. We have focused on the evolution of dust content in one of these, SBS 0335-052, since this object may be experiencing its first burst of star formation ($\\mbox{age}\\la 10^7$ yr; Vanzi et al.\\ 2000).  In this Letter, we first model the dust content of a young galaxy whose age is less than $10^8$ yr (\\S~\\ref{sec:model}). Then, we compare the model predictions with the observed quantities of the extremely low-metallicity BCD, SBS\\ 0335-052 (\\S~\\ref{sec:discussion}). Finally, we comment on implications for damped Ly$\\alpha$ systems (DLAs). ",
        "conclusions": "\\label{sec:discussion}  \\subsection{Physical Conditions of SBS 0335-052}\\label{subsec:physical}  For the dust mass, Dale et al.\\ (2001) derived the value of $M_{\\rm d}=2400~M_\\odot$ from black-body fitting of the observed spectrum of SBS 0335-052. Hunt et al.\\ (2001) derived an upper limit of $10^5~M_\\odot$ from the extinction. A starburst of age $3\\times 10^6$--$5\\times 10^7$ yr reproduces those values (Fig.\\ \\ref{fig:dust_ev}). This range is consistent with the age observationally determined by Vanzi et al.\\ (2000). For the dust temperature, it is necessary that the radius of the star-forming region be smaller than 100 pc to explain the 80 K component revealed by {\\it ISO} (Dale et al.\\ 2001), although the temperature may be overestimated because of the contamination of the emission from very small grains heated stochastically.  As observational information on the H$_2$ abundance, near-infrared emission lines are available. The line intensity is dependent on the excitation mechanism and thus the quantity of H$_2$ is poorly constrained. However, the detection of the H$_2$ lines (Vanzi et al.\\ 2000) implies that there is an effective mechanism that shields the UV-dissociative photons. The compactness as observed by Hunt et al.\\ (2001) and Dale et al.\\ (2001) may explain the existence of \\H2 for this object (Fig.\\ \\ref{fig:H2}) because of the shielding effects.  Finally, we conclude that if the star forming region of SBS 0335-052 is compact ($r_{\\rm SF}\\la 100$ pc) and the gas density is large ($\\nH\\ga 10^3$ cm$^{-3}$), the young age, the dust amount, the FIR luminosity, dust temperature, and the \\H2 abundance are mutually consistent and explained by our models.  \\subsection{Implications for High-Redshift Galaxies}  The main conclusion drawn in this Letter is that star formation activity in extremely metal-poor systems can suffer from extinction. Thus, FIR observations of primeval galaxies are important to observe such dust-enshrouded star formation activity.  Our results for the \\H2 abundance are interesting in connection with the recent observations of DLAs. There are several DLAs whose abundances of dust and H$_2$ are observationally derived. Levshakov et al.\\ (2000) assumed equilibrium between the photodissociation of H$_2$ by the background UV radiation and formation on grains. They suggested that the small dust number density (as low as 10$^{-3}$ times the Galactic value) can account for the observed small molecular fraction ($f_{\\rm H_2}\\sim 4\\times 10^{-8}$). We have shown that such a low molecular fraction can be due to the dissociation caused by recently formed OB stars. We also have suggested that if the region is compact and dense, the dust produced by SNe II can shield the UV radiation. This can produce a strong correlation between dust amount and \\H2 fraction. Indeed, a correlation between the two quantities has been observationally verified. Some DLAs whose depletion is large show molecular fractions up to 0.2 (Ge, Bechtold, \\& Kulkarni 2001). Such high values are possible when dust shielding of the UV radiation field is effective."
    },
    "0201/astro-ph0201324_arXiv.txt": {
        "abstract": "Using high resolution near-infrared spectroscopy with the Keck telescope, we have detected the radial velocity signatures of the cool secondary components in four optically identified pre$-$main-sequence, single-lined spectroscopic binaries. All are weak-lined T Tauri stars with well-defined center of mass velocities. The mass ratio for one young binary, NTTS 160905$-$1859, is M$_2$/M$_1$ $=$ 0.18$\\pm$0.01, the smallest yet measured dynamically for a pre$-$main-sequence spectroscopic binary. These new results demonstrate the power of infrared spectroscopy for the dynamical identification of cool secondaries. Visible light spectroscopy, to date, has not revealed any pre$-$main-sequence secondary stars with masses $<$0.5 M$_{\\odot}$, while two of the young systems reported here are in that range. We compare our targets with a compilation of the published young double-lined spectroscopic binaries and discuss our unique contribution to this sample. ",
        "introduction": "The most important characteristic of a star is its mass. Although there have been recent advances in determining the masses of 1$-$2 M$_{\\odot}$ pre$-$main-sequence (PMS) stars, there exist only a few measurements of the masses of sub-solar mass PMS stars \\citep{sim00b}.  As a result, theoretical PMS evolutionary tracks have not been well calibrated at the low-mass end where uncertainties in the models give rise to scatter as large as a factor of 3 in mass and 10 in age \\citep{sim00}. Dynamical measurements of the masses of close binary systems comprised of coeval components with a small mass ratio provide an initial step towards tests of the PMS tracks over a broad parameter range.  In the last decade, much progress has been made towards determining the masses of PMS stars. \\citet{ghe95}, \\citet{sim96}, and \\citet{ste01} take advantage of near-infrared (IR) speckle interferometry and the fine guidance sensors available on the {\\it HST} to map the orbits of young binary systems. Using optical wavelength ($\\sim$5000 \\AA) spectroscopy, \\citet{mat89} and \\citet{mat94} identified twelve young, low-mass, double-lined spectroscopic binaries (SB2s) and thirteen single-lined systems (SB1s). Following this seminal work, other groups have recently detected new PMS spectroscopic binary systems.  We discuss the sample of published SB2s in \\S 5.  Our immediate goal is to measure mass ratios and to enlarge the sample to study the mass ratio distribution of young stars formed in diverse environments. Precise measurements of the mass ratio distribution among young star spectroscopic systems provide critical input for theories of star formation.  For example, \\citet{bat00} shows how the density profiles of progenitor molecular cloud cores determine the distribution of mass ratios in the close ($<$ 1 AU) binary population which forms there.  Absolute masses can be obtained from SB2s which are eclipsing or from SB2s with known orbital elements determined from high resolution mapping of the binary orbits (e.g., Boden et al. 1999).  SB2s identified by optical wavelength observations tend to have mass ratios close to one.  As first described in \\citet{pra98} and further developed in \\citet{maz01}, our approach to detecting low mass ratio systems is based on the strategy that, at IR wavelengths, the flux of the secondary star is a greater fraction of the total flux than at optical wavelengths. For example, for two blackbodies of temperatures 5000 and 3000 K, the ratio of the Planck functions at 5000 \\AA~ is about 50, whereas at 1.555 $\\mu$m, the ratio is about 4. \\citet{maz01} demonstrate the efficacy of this technique, applied to three main-sequence SB1s.  \\citet{maz97} used an intermediate approach, detecting main-sequence binaries with mass ratios of $0.57\\pm0.02$ and $0.48\\pm0.03$ in the $R$ band. With the advent of powerful cross-dispersed echelle spectrometers, such as NIRSPEC on the Keck telescope, the IR approach has become even more efficient.  In this paper we describe new observations of four PMS SB1s for which we have detected the low mass companion by cross-correlation against an extensive new library of spectral type standard star templates. The observations and data reduction methods are described in \\S 2, the template spectra are presented in \\S 3 and the derived dynamical mass ratios are discussed in \\S 4.  Section 5 provides a discussion of the results, including a comparison with the complete sample of previously published SB2s and a comparison of our mass ratios with the predictions of evolutionary tracks in the H$-$R diagram.  A summary appears in \\S 6. Previous reports on applications of this technique have appeared in \\citet{pra98}, \\citet{ste01}, and \\citet{maz01}. ",
        "conclusions": "To illustrate the sensitivity of the IR technique, Figure 4 combines our sample with the IR detected SB2 NTTS 045251+3016 \\citep{ste01} and 26 optical wavelength detected SB2s.  From Figure 4 it is evident that the small sample presented in this paper includes the lowest mass ratios ever measured for PMS SB2s. Table 4 lists the entire sample of PMS SB2s plotted in Figure 4, including the location of the binaries, column (2), their mass ratios, periods, flux ratios, and primary masses (see below), in columns (3), (4) (5), and (6), and the references used for these data, column (7).  Figure 4 displays the well-known bias of SB2s detected in visible light toward mass ratios of one.  Our results start to fill in the low mass ratio part of the observed distribution.  A concentrated effort to observe a complete sample of SBs will be required to determine the true distribution of PMS SB2 mass ratios. It is interesting to notice (Table 4) that the five SB2s contributed so far by the IR technique span the same range of periods identified among the optically detected SB2s.  A complete sample of SB2s will therefore enable tests of theoretical predictions of the mass ratio distribution as a function of binary separation (e.g., Bate and Bonnell 1997).  High angular resolution imaging studies of young binaries reveal a broad range of component mass ratios, down to the regime of brown dwarfs, as inferred photometrically by comparing the location of the binary components in the H$-$R diagram with PMS evolutionary tracks \\citep{whi99, whi01, woi01}.  However, differences among the theoretical calculations of PMS evolution produce considerable scatter among the inferred mass ratios (e.g., Figure 7 of Woitas et al. 2001).  Dynamical measurements offer the means to determine the mass ratios and their distribution precisely and thus to test specific predictions of calculations of binary formation (e.g., Bate 2000).  Except for the eclipsing systems, indicated in Table 4, and the spatially resolved NTTS 0455251+3016 \\citep{ste01}, we cannot derive dynamical values for the component masses of the SB2s in Figure 4 because their orbital inclinations are not yet known. In order to compare our IR SB2 secondary masses with their minimum values set by the mass function, and to learn the mass range currently sampled by dynamical techniques, we estimate the primary masses, and thus M$_2$ using the dynamical mass ratios (Table 5).  The procedure we use to obtain M$_1$ is described in detail in Appendix A.  The M$_1$ values for the entire SB2 sample are included in Table 4, column (6).  For the 4 IR SB2s described in this paper, columns (3), (4), and (7) of Table 5 provide estimates for M$_1$, M$_2$, and the minimum mass, M$_{2_{min}}$, derived by setting the orbital inclination to $\\pi/2$ in the mass function f(M) of \\citet{mat94}, column (6). The uncertainties in M$_2$ and M$_{2_{min}}$ include the uncertainties in the mass ratio and in our estimate of M$_1$. In all cases, M$_2$ is consistently greater than its predicted minimum value.  In Figure 5 we plot mass ratio as a function of estimated M$_1$. Figure 5 shows that smaller secondary masses, as well as smaller mass ratios, are more readily identified with IR spectroscopy;  all the companions detected by the IR technique have M$_2$ $<$ 1 M$_{\\odot}$. In general, few primary star masses less than 1 M$_{\\odot}$ and no primary star masses less than 0.7 M$_{\\odot}$ have been detected in PMS SB2s.  This is illustrated by the gap on the left hand side of Figure 5 and is probably the result of a selection effect for brighter systems among SB surveys. The total sample shown in Figure 5 is weighted towards PMS stars with little or no circumstellar material, although spanning a wide range of ages, from $\\sim$ 10$^5$ yr through $\\sim$ 10$^7$ yr.  Young stars which lack circumstellar envelopes and disks do not experience veiling of their photospheric absorption lines, and hence provide easier targets for the identification of spectroscopic binaries.  The essential result of TODCOR analysis is the mass ratio of the components.  This value is robust because it is determined by cross-correlation of the many spectral lines in common between the target and templates.  The analysis also provides estimates of the primary and secondary spectral types and their $H$ band flux ratio, which give the maximum correlation with the target.  The templates almost certainly do not match exactly to the spectra of the stars in the target binaries because the surface gravities of PMS stars are lower than those of main-sequence stars, and because the templates and the PMS targets may differ in metallicity. Similarly, $H$ band flux ratios do not necessarily represent accurate component luminosity ratios. Nonetheless, it is of interest to investigate the extent to which the values of $M_2/M_1$ implied by our dynamical results are consistent with PMS evolutionary tracks in the H$-$R diagram, presented in Figure 6.  Our approach to estimating the effective temperature, T$_{eff}$, and luminosity for each object, and the associated uncertainties, is described in Appendix B. We use the tracks of Palla \\& Stahler (1997) because they are representative of PMS evolution, similar to the results of Baraffe et al. (1998) and Siess et al. (2000), and span the range of masses of our sample objects. Figure 6 shows that, to within 1 $\\sigma$, the secondary masses implied by the mass ratios (Table 5) are consistent with the secondary masses estimated from the tracks (column (5) of Table 5), except for Parenago 1925. Only for NTTS 160905$-$1859 do the components lie on distinct isochrones; the other systems appear to be coeval to 1 $\\sigma$. It will be possible to derive absolute, dynamical masses for these objects within the next decade as ground and space based interferometers become available for mapping their orbits. These mass determinations will be precise to within a few percent or less, a requirement for meaningful testing of models. Reliable luminosity estimates, however, cannot be obtained by dynamical means and will require not only spatially resolved observations, but also multi-wavelength data in order to obtain precise values."
    },
    "0201/physics0201069_arXiv.txt": {
        "abstract": "Mark  correlations provide a  systematic approach  to look  at objects both  distributed  in space  and  bearing  intrinsic information,  for instance  on  physical  properties.   The interplay  of  the  objects' properties (marks)  with the spatial  clustering is of  vivid interest for many applications; are, e.g., galaxies with high luminosities more strongly clustered than dim ones?   Do neighbored pores in a sandstone have similar sizes?  How does the  shape of impact craters on a planet depend on the geological surface properties?  In this article, we give an introduction  into the  appropriate mathematical framework  to deal with  such questions,  i.e.   the theory  of  marked point  processes. After having  clarified the notion  of segregation effects,  we define universal test quantities applicable to realizations of a marked point processes.  We show their power  using concrete data sets in analyzing the luminosity-dependence  of the galaxy clustering,  the alignment of dark   matter  halos  in   gravitational  $N$-body   simulations,  the morphology- and diameter-dependence of the Martian crater distribution and the size correlations of pores in sandstone. In order to understand our data in more detail, we discuss the Boolean depletion  model, the  random field  model  and the  Cox random  field model.    The  first   model  describes   depletion  effects   in  the distribution of  Martian craters and  pores in sandstone,  whereas the last   one  accounts   at   least  qualitatively   for  the   observed luminosity-dependence of the galaxy clustering. ",
        "introduction": "\\label{sec:kerscher_basic} Observations of  spatial patterns at various  length scales frequently are the only point where  the physical world meets theoretical models. In many cases these patterns consist of a number of comparable objects distributed in space  such as pores in a sandstone,  or craters on the surface    of   a    planet.     Another   example    is   given    in Figure~\\ref{fig:kerscher_galaxies-circles},   where  we   display   the  galaxy distribution as traced by a  recent galaxy catalogue. The galaxies are represented as  circles centered at their positions,  whereas the size of the circles mirrors the luminosity of a galaxy. In order to test to which  extent  theoretical   predictions  fit  the  empirically  found structures  of that  type, one  has to  rely on  quantitative measures describing the physical  information.  Since theoretical models mostly do not try to explain  the structures individually, but rather predict some of their generic properties,  one has to adopt a {\\em statistical point of view} and to interpret  the data as a realization of a random process.  In  a first step one  often confines oneself  to the spatial distribution of the objects constituting the patterns and investigates their clustering  thereby thinking  of it as  a realization of  a {\\em point process}.   Assuming that  perspective, however, one  neglects a possible  linkage between  the  spatial clustering  and the  intrinsic properties of the objects.  For instance, there are strong indications that the clustering of galaxies depends on their luminosity as well as on        their         morphological        type.         Considering Figure~\\ref{fig:kerscher_galaxies-circles},  one   might  infer  that  luminous galaxies are more strongly correlated than dim ones. Effects like that are referred to as {\\em mark segregation} and provide insight into the generation and interactions of,  e.g., galaxies or other objects under consideration. The  appropriate statistical framework  to describe the relation between the spatial distribution of \\begin{figure} \\begin{center} \\epsfig{file=kerscher_figs/ssrs_circ.eps,width=8cm} \\end{center} \\caption{The  galaxy  distribution  as  traced  by  the  Southern  Sky Redshift Survey 2 (SSRS~2). We show a part of the sample investigated, projected down into two  dimensions.  Each circle represents a galaxy, its  radius is proportional  to the  galaxy's luminosity.  For further details see Section~\\ref{sec:kerscher_gal}.\\label{fig:kerscher_galaxies-circles}} \\end{figure} physical objects and their inner properties are {\\em marked point processes}, where discrete, scalar-, or vector-valued marks are attached to the random points.\\\\ In this contribution we outline how to describe marked point processes; along that line we discuss two notions of independence (Section~\\ref{sec:kerscher_basic}) and define corresponding statistics that allow us to quantify possible dependencies.  After having shown that some empirical data sets show significant signals of mark segregation (Section{}\\ref{sec:kerscher_data}), we turn to analytical models, both motivated by mathematical and physical considerations (Section~{}\\ref{sec:kerscher_models}). \\\\ Contact  distribution functions  as presented  in the  contribution by D.~Hug et al.  in this volume are an  alternative technique to measure and  statistically quantify distances  which finally  can be  used to relate  physical properties  to spatial  structures.  Mark correlation functions  are useful  to  quantify molecular  orientations in  liquid crystals (see  the contribution by  F.~Schmid and N.~H.~Phuong  in this volume)   or   in  self-assembling   amphiphilic   systems  (see   the contribution by U.~S.~Schwarz and G.~Gompper in this volume). But also to  study anisotropies  in composite  or porous  materials,  which are essential for elastic and transport properties (see the contributions by D.~Jeulin,  C.~Arns et  al. and H.-J.~Vogel  in this  volume), mark correlations may be relevant. \\subsection{The framework} The empirical data -- the positions ${\\mathbf{x}}_i$ of some objects together with their intrinsic properties $m_i$ -- are interpreted as a realization of a marked point process $\\{({\\mathbf{x}}_i,m_i)\\}_{i=1}^N$ (Stoyan, Kendall and Mecke, 1995). For simplicity we restrict ourselves to homogeneous and isotropic processes. \\\\ The hierarchy of joint probability densities provides a suitable tool to describe the stochastic properties of a marked point process. Thus, let $\\varrho_1^{{\\cal S}{\\cal M}}\\left(({\\mathbf{x}},m)\\right)$ denote the probability density of finding a point at ${\\mathbf{x}}$ with a mark $m$. For a homogeneous process this splits into $\\varrho_1^{{\\cal S}{\\cal M}}\\left(({\\mathbf{x}},m)\\right)=\\varrho{\\cal M}_1(m)$ where $\\varrho$ denotes the mean number density of points in space and ${\\cal M}_1(m)$ is the probability density of finding the mark $m$ on an arbitrary point. Later on we need moments of this mark distribution;  for real-valued marks the $k$th-moment of the mark-distribution is defined as \\begin{equation} \\overline{m^k} = \\int{\\rm d} m\\ {\\cal M}_1(m) m^k; \\end{equation} the mark variance is $\\sigma_M^2=\\overline{m^2}-\\overline{m}^2$. \\\\ Accordingly,   $\\varrho_2^{{\\cal S}{\\cal M}}\\left(({\\mathbf{x}}_1,m_1),({\\mathbf{x}}_2,m_2)\\right)$ quantifies the probability  density to find two points  at ${\\mathbf{x}}_1$ and ${\\mathbf{x}}_2$  with marks  $m_1$ and  $m_2$, respectively  (for second-order theory        of        marked        point       processes        see {}\\cite{kerscher_stoyan:stochgeom,kerscher_stoyan:fractals}).  It effectively depends only on $m_1$, $m_2$, and  the pair separation $r=|{\\mathbf{x}}_2-{\\mathbf{x}}_1|$ for a homogeneous and isotropic  process. Two-point properties certainly are the simplest non-trivial  quantities for homogeneous random processes, but it may be necessary to  move on to higher correlations in order to discriminate between certain models.  \\subsection{Two notions of independence} In the following we will discuss two notions of independence, which may arise for marked point patterns. For this, consider two Renaissance families, call them the Sforza and the Gonzaga.  They used to build castles spread out more or less homogeneously over Italy. In order to describe this example in terms of a marked point process, we consider the locations of the castles as points on a map of Italy, and treat a castle's owner as a discrete mark, $S$ and $G$, respectively.  There are many ways how the castles can be built and related to each other.   \\paragraph{Independent sub-point processes:} For example,  the Sforza  may build their  castles regardless  of the Gonzaga  castles. In  that case  the probability  of finding  a Sforza castle at ${\\mathbf{x}}_1$ and a  Gonzaga castle at ${\\mathbf{x}}_2$ factorizes into two one-point  probabilities and  we  can  think of  the  Sforza and  the Gonzaga castles as uncorrelated  sub-point processes. In the language of marked point processes this means, e.g., that \\begin{gather} \\begin{split} \\varrho_2^{{\\cal S},{\\cal M}}\\left(  ({\\mathbf{x}}_1,m_1)  ,({\\mathbf{x}}_2,m_2)\\right) &= \\varrho_1^{{\\cal S}{\\cal M}}\\left(  ({\\mathbf{x}}_1,m_1)\\right)\\varrho_1^{{\\cal S}{\\cal M}}\\left( ({\\mathbf{x}}_2,m_2)\\right) \\\\ &= \\varrho^2 {\\cal M}_1(m_1){\\cal M}_1(m_2) , \\end{split} \\end{gather} for any $m_1\\neq m_2$. If all the joint $n$-point densities factorize into a product of $n'$-point densities of one type each, then we speak of {\\em independent sub-point processes}.  Dependent sub-point processes indicate {\\em interactions} between points of different marks; for instance, the Gonzaga may build their castles close to the Sforza ones in order to avoid that a region becomes dominated by the other family's castles.  \\paragraph{Mark-independent clustering:} A second type of independence  refers to  the question  whether the  different families have  different  styles to  plan  their  castles.   For instance,  the Gonzaga  may distribute  their  castles in  a  grid-like manner  over Italy, whereas the  Sforza may incline to build  a second castle close to each  castle they  own. Rather than  asking whether  two sub-point processes (namely  the Gonzaga  and the Sforza  castles, respectively) are  independent  (``independent   sub-point  processes''),  we  are  now discussing  whether  they  are   {\\em  different}  as  regards  their statistical clustering properties. Any such difference means that the clustering {\\em depends}  on the  intrinsic mark of  a point. \\\\ Whenever the two-point probability  density of finding two objects  at ${\\mathbf{x}}_1$ and ${\\mathbf{x}}_2$ depends on the objects'  intrinsic properties we speak of {\\em mark-dependent clustering}.  It is  useful to rephrase this statement by using Bayes' theorem and the conditional mark probability density \\begin{gather} \\label{eq:kerscher_cond-mark-density} {\\cal M}_2(m_1,m_2|{\\mathbf{x}}_1,{\\mathbf{x}}_2)= \\frac{\\varrho_2^{{\\cal S},{\\cal M}}\\left(({\\mathbf{x}}_1,m_1),({\\mathbf{x}}_2,m_2)\\right)} {\\varrho_2^{{\\cal S}}\\left({\\mathbf{x}}_1,{\\mathbf{x}}_2\\right)}, \\end{gather} in case the spatial product density $\\varrho_2^{{\\cal S}}(\\cdot)$ does not vanish.  ${\\cal M}_2(m_1,m_2|{\\mathbf{x}}_1,{\\mathbf{x}}_2)$ is the probability density of finding the marks $m_1$ and $m_2$ on objects located at ${\\mathbf{x}}_1$ and ${\\mathbf{x}}_2$, given that there are objects at these points.  Clearly, ${\\cal M}_2(m_1,m_2|{\\mathbf{x}}_1,{\\mathbf{x}}_2)$ depends only on the pair separation $r=|{\\mathbf{x}}_1-{\\mathbf{x}}_2|$ for homogeneous and isotropic point processes.  We speak of {\\em mark-independent} clustering, if ${\\cal M}_2(m_1,m_2|r)$ factorizes \\begin{equation} {\\cal M}_2(m_1,m_2|r) = {\\cal M}_1(m_1) {\\cal M}_1(m_2) \\end{equation} and thus does not depend on the pair separation. That means that regarding their marks, pairs with a separation $r$ are not different from any other pairs.  On the contrary, mark-dependent clustering or {\\em mark segregation} implies that the marks on certain pairs show deviations from the global mark distribution. \\\\ In order to distinguish between both sorts of independencies, let us consider the case where we are given a map of Italy only showing the Gonzaga castles.  If the distribution of castles in Italy can be understood as consisting of independent sub-point processes, we cannot infer anything about the Sforza castles from the Gonzaga ones. However, if $\\varrho_2^{{\\cal S},{\\cal M}}\\left(({\\mathbf{x}}_1,S),({\\mathbf{x}}_2,G)\\right) >\\varrho^2 {\\cal M}_1(S){\\cal M}_1(G)$, Sforza castles are likely to be found close to Gonzaga ones. Here, ${\\cal M}_1(S)$ and ${\\cal M}_1(G)$ are the probabilities that a castle belongs to the Sforza or Gonzaga family. If, on  the other  hand, mark-independent clustering  applies, typical clustering  properties such  as  the spatial  clustering strength  are equal  for both  castle  distributions, and  the  Gonzaga castles  are in the statistical sense already  representative of  the  whole castle  distribution in  Italy. That means in  particular that, if the Gonzaga  castles are clustered, so are the Sforza ones.  Before we turn to applications, we have to develop practical test quantities in order to test for segregation effects in real data and to describe them in more detail.   \\subsection{Investigating the independence of sub-point processes}   To  investigate  correlations  between sub-point  processes,  suitably extended nearest neighbor distribution functions or $K$-functions have been   employed   {}\\cite{kerscher_cox:point,kerscher_diggle:statistical}.   Also   the (conditional)   cross-correlation   functions   can   be   used   (see Eq.~\\ref{eq:kerscher_cond-crosscorr}), for a further test see {}\\cite{kerscher_stoyan:fractals}, p.~302. Here we consider a multivariate extension of the $J$-function {}\\cite{kerscher_vanlieshout:j}, as suggested by {}\\cite{kerscher_vanlieshout:indices}. \\\\ For this, consider the nearest neighbor's distance distribution from an object with mark $m_i$ to other objects with mark $m_j$, $G_{ij}(r)$ (``$i$ to $j$'', for details see {}\\cite{kerscher_vanlieshout:indices}).  Let $G_{i\\circ}(r)$ denote the distribution of the nearest neighbor's distance from an object of type $i$ to any other object (denoted by $\\circ$).  Finally, $G_{\\circ\\circ}(r)$ is the nearest neighbor distribution of all points.  Similar extensions of the empty space function are possible, too. Let $F_i(r)$ denote the distribution of the nearest $i$-object's distance from an arbitrary position, whereas $F_\\circ(r)$ is the nearest object's distance distribution from a random point in space to any object in the sample.  We consider the following quantities: \\begin{gather} \\label{eq:kerscher_def-multivar-J} \\begin{split} J_{ij}(r) =\\frac{1-G_{ij}(r)}{1-F_j(r)} ,\\ J_{i\\circ}(r) =\\frac{1-G_{i\\circ}(r)}{1-F_\\circ(r)} ,\\ J(r) =\\frac{1-G_{\\circ\\circ}(r)}{1-F_\\circ(r)} , \\end{split} \\end{gather} They are defined whenever $F_j(r),F_\\circ(r) <1$. If two sub-point processes, defined by marks $i\\neq j$, are independent then one gets {}\\cite{kerscher_vanlieshout:indices} \\begin{gather} \\label{eq:kerscher_jijone} J_{ij}(r) = 1 . \\end{gather} Note, that the $J_{ij}$ depend on higher-order correlations functions, similar to the $J$-function {}\\cite{kerscher_kerscher:constructing}.  Suitable estimators for these $J$-functions are derived from estimators of the $F$ and $G$-functions {}\\cite{kerscher_stoyan:stochgeom,kerscher_baddeley:sampling}.  \\subsection{Investigating mark segregation} \\label{sec:kerscher_mark-segregation} In order to quantify the  mark-dependent clustering or to look for the mark  segregation,  it  proves  useful to  integrate  the  conditional probability density ${\\cal M}_2(m_1,m_2|r)$ over the marks weighting with a test                       function                       $f(m_1,m_2)$ {}\\cite{kerscher_stoyan:oncorrelations,kerscher_stoyan:stochgeom}.     This   procedure reduces the number  of variables and leaves us  with the weighted pair average: \\begin{equation} \\label{eq:kerscher_def-paverage} \\paverage{f} = \\int{\\rm d}{m_1}\\int{\\rm d}{m_2}\\ f(m_1,m_2) {\\cal M}_2(m_1,m_2|r) . \\end{equation} The choice of an appropriate weight-function depends on whether the marks are non-quantitative labels or continuous physical quantities. \\begin{enumerate} \\item[1.\\ ] For labels only combinations of indicator-functions are possible,  the integral  degenerates into a  sum over  the labels. Supposed  the marks  of our  objects belong  to classes  labelled with $i,j,\\ldots$, the conditional cross-correlation functions are given by \\begin{equation} \\label{eq:kerscher_cond-crosscorr} C_{i j}(r) \\equiv \\paverage{\\delta_{m_1i} \\delta_{m_2j} + (1-\\delta_{ij}) \\delta_{m_2i} \\delta_{m_1j}} (r), \\end{equation} with  the Kronecker $\\delta_{m_1i}=1$  for $m_1=i$  and zero otherwise. Mark           segregation          is           indicated          by $C_{ij}\\ne2\\varrho_i\\varrho_j/\\varrho^2$    for     $i\\ne    j$    and $C_{ii}\\ne\\varrho_i^2/\\varrho^2$, where $\\varrho_i$ denotes the number density of points with  label $i$.  The $C_{ij}$ are cross-correlation functions under the {\\em condition} that two points are separated by a distance  of~$r$  (compare  {}\\cite{kerscher_stoyan:fractals},  p.~264,  for applications   see  the   Martian  crater   distribution   studied  in Sect.~\\ref{sec:kerscher_martian}    and    Figure~\\ref{fig:kerscher_mars-markcorr}    in particular). \\item[2.\\  ] For positive  real-valued marks  $m$, the  following pair averages     prove     to      be     powerful     and     distinctive {}\\cite{kerscher_schlather:mark,kerscher_beisbart:luminosity}: \\\\ \\noindent \\begin{enumerate} \\item One of the most simplest weights to be used is the mean mark: \\begin{equation} k_{m}(r) \\equiv \\frac{\\paverage{m_1+m_2}(r)}{2\\ \\overline{m}} . \\end{equation} quantifies the deviation of the mean mark on pairs with separation $r$ from the overall mean mark $\\overline{m}$.  A $k_{m}>1$ indicates mark segregation for point pairs with a separation $r$, specifically their mean mark is then larger than the overall mark average.\\\\ Closely related is Stoyan's $k_{mm}$ function using the squared geometric mean of the marks as a weight {}\\cite{kerscher_stoyan:oncorrelations,kerscher_stoyan:fractals} \\begin{equation} k_{mm}(r) \\equiv \\frac{\\paverage{m_1m_2}(r)}{\\overline{m}^2} . \\end{equation} \\item Accordingly, higher moments of the marks may be used to quantify mark segregation, like the mark fluctuations \\begin{equation} \\var(r) \\equiv \\paverage{\\left(m_1-\\paverage{m_1}(r)\\right)^2}(r) , \\end{equation} or the mark-variogram {}\\cite{kerscher_waelder:variograms,kerscher_stoyan:variograms}: \\begin{equation} \\gamma(r) \\equiv \\paverage{\\tfrac{1}{2}\\left(m_1-m_2\\right)^2}(r) , \\end{equation} \\item The mark covariance {}\\cite{kerscher_cressie:statistics} is \\begin{equation} \\label{eq:kerscher_def-cov} \\cov(r) \\equiv \\paverage{m_1 m_2}(r) -\\paverage{m_1}(r)\\paverage{m_2}(r) . \\end{equation} Mark segregation can be detected by looking whether $\\cov(r)$ differs from zero. A $\\cov(r)$ larger than zero, e.g., indicates that points with separation $r$ tend to have similar marks.  Sometimes the mark covariance is normalized by the fluctuations {}\\cite{kerscher_isham:marked}: $\\cov(r)/\\var(r)$. \\end{enumerate} These conditional mark correlation functions can be calculated from only three independent pair averages~\\cite{kerscher_schlather:mark}: $\\paverage{m}(r)$, $\\paverage{m_1m_2}(r)$, and $\\paverage{m^2}(r)$. Thus the above mentioned characteristics are not independent, e.g.\\ $\\var(r)=\\gamma(r)+cov(r)$. \\\\ We apply  these mark correlation functions to  the galaxy distribution in   Section~\\ref{sec:kerscher_gal}  (Figure~\\ref{fig:kerscher_ssrs-lum}),   to  Martian craters  in Section~\\ref{sec:kerscher_martian} (Figure~\\ref{fig:kerscher_mars-markcorr}) and to pores in sandstones considered in Section~\\ref{sec:kerscher_sandstone}. \\item[3.\\ ] Also vector-valued information ${\\mathbf{l}}_i$, describing, e.g., the orientation of an anisotropic object at  position ${\\mathbf{x}}_i$ may be available. It is therefore interesting to consider vector marks such as done  by {}\\cite{kerscher_ohser:onsecond,kerscher_penttinen:statistical,kerscher_stoyan:fractals} who use a mark correlation function to quantify the alignment of vector marks. Here   we  suggest  three   mark  correlation   functions  quantifying geometrically different  possibilities of  an alignment.  In  order to ensure coordinate-independence of our  descriptors, we focus on scalar combinations of the  vector marks in using the  scalar product $\\cdot$ and the  cross product  $\\times$.  Different from  the case  of scalar marks,  it  is  a  non-trivial  task  to find  a  set  of  vector-mark correlation functions which contain all possible information (at least up to a  fixed order in mark space).  We  provide a systematic account of how  to construct suitable  vector-mark correlation functions  in a complete and unique way for general dimensions in the Appendix. \\\\ Here we only cite  the most important  results.  For that we need  the distance vector   between   two  points,  ${\\mathbf{r}}\\equiv{\\mathbf{x}}_1-{\\mathbf{x}}_2$,  the normalized distance vector, $\\hat{\\mathbf{r}}\\equiv{\\mathbf{r}}/r$,  and the normalized  vector mark: $\\hat{\\mathbf{l}}_i\\equiv{\\mathbf{l}}_i/l_i$  with $l_i=|{\\mathbf{l}}_i|$. The    following conditional mark   correlation  functions will  be  used  to  quantify alignment effects: \\begin{enumerate} \\item ${\\cal A}(r)$ quantifies the ${\\cal A}$lignment of the two vector marks ${\\mathbf{l}}_1$ and ${\\mathbf{l}}_2$: \\begin{equation} {\\cal A}(r) = \\frac{1}{\\overline{l}^2} \\paverage{{\\mathbf{l}}_1\\cdot{\\mathbf{l}}_2}(r)\\;\\;\\;. \\label{eq:kerscher_align} \\end{equation} It is proportional to the cosine of the angle between ${\\mathbf{l}}_1$ and ${\\mathbf{l}}_2$. We normalize  with the mean $\\overline{l}$. For purely independent vector marks ${\\cal A}(r)$ is zero, whereas ${\\cal A}(r)>0$ means that the marks of pairs separated by $r$ tend to align parallel to each other. -- In some applications, e.g. for the orientations of ellipsoidal objects, the vector mark is only defined up to a sign,  i.e.\\ ${\\mathbf{l}}$ and $-{\\mathbf{l}}$ mean actually the same. In this case the absolute value of the scalar product is useful: \\begin{equation} \\label{eq:kerscher_def-vector-corr-absolute} {\\cal A}'(r) \\equiv \\frac{1}{\\overline{l}^2} \\paverage{|{\\mathbf{l}}_1\\cdot{\\mathbf{l}}_2|}(r)\\;\\;\\;. \\end{equation} For uncorrelated random vectors we get ${\\cal A}'(r)=1/2$. ${\\cal A}$ and ${\\cal A}'$ can  readily be  generalized to  any  dimension $d$,  where we  expect ${\\cal A}'=$ $\\pi^{-\\frac{1}{2}}\\frac{\\Gamma(\\frac{d}{2})}{\\Gamma(\\frac{d+1}{2})}$ for  uncorrelated random  orientations.  In two  dimensions ${\\cal A}'$  is proportional to $k_d$ as defined by~\\cite{kerscher_stoyan:fractals}. \\item ${\\cal F}(r)$  quantifies  the  ${\\cal F}$ilamentary  alignment of  the  vectors ${\\mathbf{l}}_1$  and ${\\mathbf{l}}_2$  with respect  to the  line connecting  both halo positions: \\begin{equation} {\\cal F}(r) \\equiv \\frac{1}{2\\ \\overline{l}} \\paverage{|{\\mathbf{l}}_1\\cdot{\\hat{\\mathbf{r}}}| + |{\\mathbf{l}}_2\\cdot{\\hat{\\mathbf{r}}}|}(r), \\label{eq:kerscher_filament} \\end{equation} ${\\cal F}(r)$ is  proportional to the  cosine of the angle  between ${\\mathbf{l}}_1$ and  the  distance  vector   $\\hat{\\mathbf{r}}$  connecting  the  points.   For uncorrelated  random  vector  marks,  we  expect  again  ${\\cal F}(r)=1/2$; ${\\cal F}(r)$ becomes  larger than that,  whenever the vector marks  of the objects tend  to point to  objects separated by  $r$ -- an  example is provided  by  rod-like metallic  grains  in  an  electric field:  they concentrate along  the field lines  and orient themselves  parallel to the field lines. \\item ${\\cal P}(r)$ quantifies the ${\\cal P}$lanar alignment of the vectors and the distance vector. ${\\cal P}(r)$ is proportional to the volume of the rhomb defined by ${\\mathbf{l}}_1$, ${\\mathbf{l}}_2$ and $\\hat{\\mathbf{r}}$: \\begin{equation} {\\cal P}(r) = \\frac{1}{2\\overline{l}^2}\\ \\paverage{\\left| {\\mathbf{l}}_1\\cdot\\frac{{\\mathbf{l}}_2\\times{\\hat{\\mathbf{r}}}}{|\\hat{\\mathbf{l}}_2\\times{\\hat{\\mathbf{r}}}|} \\right| + \\left| {\\mathbf{l}}_2\\cdot\\frac{{\\mathbf{l}}_1\\times{\\hat{\\mathbf{r}}}}{|\\hat{\\mathbf{l}}_1\\times{\\hat{\\mathbf{r}}}|} \\right|}(r), \\label{eq:kerscher_planar} \\end{equation} Quite  obviously, this quantity  can not  be generalized  to arbitrary dimensions;  the deeper  reason  for  that will  become  clear in  the Appendix.   --  We get  ${\\cal P}(r)=1/2$  for  randomly oriented  vectors, whereas  it  is   becoming  larger  for  the  case   that  ${\\mathbf{l}}_2$  is perpendicular to ${\\mathbf{l}}_1$ as well as to $\\hat{\\mathbf{r}}$. \\end{enumerate} Applications of vector marks  can be found in Section~\\ref{sec:kerscher_halos} (Figure~\\ref{fig:kerscher_halos-orientation}) where we consider the orientation of dark matter halos in cosmological simulations. But one can think of other   applications:  mark   correlation  functions   may   serve  as orientational  order  parameters  in   liquid  crystals  in  order  to discriminate between nemetic and  smectic phases (see the contribution by F.~Schmid and N.~H.~Phuong in  this volume). They can also quantify the  local orientation  and  order  in liquids  such  as the  recently measured   five-fold    local   symmetry   found    in   liquid   lead {}\\cite{kerscher_reichert:lead}.  As a further application one could try to measure the signature of hexatic phases in two-dimensional colloidal dispersions and in 2D melting scenarios occurring in experiments and simulations of hard-disk systems (for a review on hard sphere models see~\\cite{kerscher_loewen:lnp}. Finally,  the orientations   of   anisotropic  channels   in   sandstone  (see   the contribution  by  C.~Arns et  al.  in  this  volume) are  relevant  for macroscopic   transport  properties,   therefore   their  quantitative characterization  in  terms of  mark  correlation  functions might  be interesting.  \\end{enumerate} Before we move on to applications  a few general remarks are in order: First,  the definition  of  these mark  characteristics  based on  the conditional density  ${\\cal M}_2(\\cdot)$ leads to ambiguities  at $r$ equal zero as discussed by {}\\cite{kerscher_schlather:mark}, but there is no problem for $r>0$. -- Furthermore, suitable estimators for  our test quantities are based on   estimators  for   the  usual   two-point   correlation  function {}\\cite{kerscher_stoyan:fractals,kerscher_capobianco:autocovariance,kerscher_beisbart:luminosity}. \\\\ Mark-dependent  clustering  can also  be  defined  at any  $n$-point level.   Mark-in\\-de\\-pendent clustering  at  every order  is called  the random  labelling   property  {}\\cite{kerscher_cox:point}.   Mark  correlation functions based on the $n$-point densities may be used.  For discrete marks        the        multivariate        $J$-functions        (see Eq.~\\eqref{eq:kerscher_def-multivar-J})   are   an   interesting   alternative, sensitive  to   higher-order  correlations.   The   random  labelling property then leads to the relation \\begin{gather} J_{i\\circ}(r) = J, \\end{gather} which may be used as a test {}\\cite{kerscher_vanlieshout:indices}. ",
        "conclusions": "Whenever objects are sampled together with their spatial positions and some of their intrinsic properties, marked point processes are the stochastic models for those data sets.  Combining the spatial information and the objects' inner properties one can constrain their generation mechanism and their interactions. \\\\ Developing the framework of marked point processes further and outlining some of their general notions is thus of interest for physical applications.  Let us therefore look at mark correlations again from both a statistical and a physical perspective.  We focused on two kinds of dependencies. \\\\ On the one hand, one can always ask, whether objects of different types ``know'' from each other.  From a statistical point of view, this is the question whether the marked point process consists of two completely independent sub-point processes.  Physically, this concerns the question whether the objects have been generated together and whether they interact with each other. \\\\ On the other hand, it is often interesting to know whether the spatial distribution of the objects  changes with their inner properties.  For the  statistician,  this translates  into  the  question whether  mark segregation  or  mark-independent   clustering  is  present.  For  the physicist such  a dependency is  interesting since one can  learn from them whether  and how  the interactions distinguish  between different object classes or  whether the formation of the  objects' mark depends on the environment. \\\\ We discussed statistics capable of probing to which extent mark correlations are present in a given data set, and showed how to assess the statistical significance.  Applying our statistics to real data, we could demonstrate, that the clustering of galaxies depends on their luminosities.  Large scale correlations of the orientations of dark matter halos were found.  Using the Mars data we could validate a picture of crater generation on the Martian surface: mainly, the local geological setting determines the crater type. We also could show that the sizes of pores in sandstone are correlated. \\\\ In order to understand empirical data sets in detail, we need models to compare to. As generic models the Boolean depletion model, the random field model and its extension, the Cox random field models are of interest. \\\\ Further application of the mark correlations properties may  inspire the development of  further models.  It seems therefore  that marked point processes could spark  interesting interactions between physicists and mathematicians.   Certainly,  the   distributions  of  physicists  and mathematicians  in  coffee breaks  at  the  Wuppertal conference  were clustered, each.  But could one  observe positive cross-correlations? Using mark  correlations we argue, that,  even more, there  is lots of space for positive interactions.\\ldots.   \\subsubsection*"
    },
    "0201/astro-ph0201330_arXiv.txt": {
        "abstract": "{We report observations of the low-luminosity $z = 5.50$ quasar RD~J030117+002025 (RD0301 hereafter) at 250 GHz (1.20~mm) using the Max-Planck Millimeter Bolometer (MAMBO) array at the IRAM 30-meter telescope. The quasar was detected with a 1.2~mm flux density of $0.87 \\pm 0.20 \\, \\rm mJy$. The lack of detectable 1.4 GHz radio emission indicates that the millimeter emission is of thermal nature, making RD0301 the most distant dust-emission source known. When matching a 50~K grey body thermal far-infrared (FIR) spectrum to the observed millimeter flux we imply a FIR luminosity $\\approx \\rm 4 \\times 10^{12} \\, L_{\\odot}$, which is comparable to the quasar's optical luminosity.  If the FIR luminosity arises from massive star formation, the implied star formation rate would be $\\sim 600\\,\\rm M_\\odot yr^{-1}$, comparable to that of the starburst galaxies which dominate the average star formation and FIR emission in the early Universe.  The FIR luminosity of RD0301 is close to the average of that found in optically far more luminous high-redshift quasars.  The comparably high millimeter to optical brightness ratio of RD0301 is further evidence for that there is no strong correlation between the optical and millimeter brightness of high-redshift quasars, supporting the idea that in high-redshift quasars the dust is not heated by the AGN, but by starbursts. ",
        "introduction": "Many high-redshift quasars (QSOs) have recently been found through wide field imaging surveys, notably the Sloan Digital Sky Survey (SDSS; Schneider et al. \\cite{sch01}; Anderson et al.  \\cite{and01}) and the Digital Palomar Sky Survey (DPSS; Kennefick et al. \\cite{ken95a},\\cite{ken95b}; Djorgovski et al. \\cite{djo99}). Now over 250 QSOs are known with redshifts $z>3.6$, thirteen of which are at $z>5$ (Fan et al. \\cite{fan99}, \\cite{fan00a}, \\cite{fan00b}; Zheng et al.  \\cite{zhe00}; Stern et al.  \\cite{ste00}; Sharp et al. \\cite{sha01}; Fan et al. \\cite{fan01}) and one at $z>6$ (Fan et al. \\cite{fan01}).  All but one of the $z>5$ QSOs were found from magnitude $i\\approx 21.5$ limited surveys, limiting their implied rest frame blue magnitudes to $M_{\\rm B} < -26$.  In a 74 arcmin$^2$ field observed for the SPICES survey (Stern et al. \\cite{ste01}), Stern et al. (\\cite{ste00}) discovered the redshift 5.50 QSO RD (i.e., ``R-band Dropout'') J030117+002025 (RD0301 hereafter), which has an $i$-band AB magnitude of 23.4 and $M_{\\rm B} \\approx -23.4$ in an Einstein-de Sitter Universe ($H_0=50~ \\rm km~s^{-1} Mpc^{-1}$, $q_0=0.5$), and $M_{\\rm B} \\approx -24.0$ in a $\\Lambda$-cosmology ($H_0=65\\rm~km~s^{-1}~Mpc^{-1}$, $\\Omega_\\Lambda=0.7$, $\\Omega_{\\rm m}=0.3$) which we adopt throughout this paper.  RD0301 is significantly fainter than the typical QSO found in the wide field SDSS and DPSS surveys.  The discovery of such a very high redshift QSO was surprising because the surface density of such objects, as implied by the intermediate-redshift QSO luminosity function and redshift evolution models, is so low that the probability to find one in the observed field is only $\\sim 15\\%$ (Stern et al. \\cite{ste00}).  The discovery may therefore imply that the luminosity evolution of faint high-redshift QSOs is weaker than expected, and that there could be a larger, optically faint QSO population at high redshifts. This possibility is particularly interesting in terms of which sources re-ionize the early Universe (cf., the recent Gunn-Peterson trough results of Becker et al. [\\cite{bec01}] and Djorgovski et al.  [\\cite{djo01}]).   Another important question in studies of the early Universe is how the growth of massive black holes relates to the formation of early stellar populations. A tight correlation between black hole masses and bulge luminosities or velocity dispersions in local galaxies indicates that their formation was closely related (Magorrian et al. \\cite{mag98}). The recent detection of strong thermal dust emission from many high-redshift QSOs shows that vigorous star formation is coeval with black hole growth (Carilli et al. \\cite{car01a}; Omont et al. \\cite{omo01}; Isaak et al. \\cite{isa01}).  Interestingly, the FIR luminosity of these QSOs does not seem to be well correlated with their optical luminosity in the studied range of optical luminosities, $M_{\\rm B} = -26$ to $-29.5$.  A similar result was previously found also in lower redshift, lower optical luminosity QSOs (Sanders et al. \\cite{san89}; Chini et al. \\cite{chi89}).  Such a lack of correlation might suggest that the warm dust is not heated by the QSO. Instead, star formation and black hole accretion could have had a common cause, which might be the infall of material toward the center of the host galaxy.  The QSOs observed in the millimeter and submillimeter surveys are the optically most luminous objects, and their implied optical luminosities are on average ten times higher than their FIR luminosities, so that dust heating by the AGN cannot be excluded on energetic grounds.  It is of great interest therefore to investigate whether strong millimeter emission is also found in high-redshift QSOs which are optically much fainter than those observed so far. In this respect, RD0301 is a unique object: it is the optically least luminous QSO known at $z>4$, and it is one of the highest redshift quasars ever detected.  Searching for high-redshift dust emission is of great interest also because for them there was little time to produce the observed dust. At redshift 5.5, the universe was about 1~Gyr old, a time not much longer than the dynamical timescale of typical galaxies.  Observing large amounts of dust at such redshifts thus sets back the first epochs of vigorous star formation irrespective of whether the observed dust was heated by an AGN or a starburst. ",
        "conclusions": "\\subsection{Thermal nature of the millimeter emission}  At redshift 5.5 our observing frequancy of 250 GHz corresponds to an emitted frequency of 1625~GHz or a wavelength of 185~$\\rm \\mu m$.  From a flux measurement at only one frequency we cannot distinguish whether the emission follows a steep thermal spectrum, or constitutes the millimeter extension of a radio synchrotron spectrum. The VLA 1.4~GHz FIRST survey (Becker, White \\& Helfand \\cite{bec95}) however yields a 5$\\sigma$ upper limit flux density of $\\simeq$~1~mJy within 30$^{\\prime\\prime}$ from RD0301. Since rising non-thermal spectra are not observed at high radio frequencies, it is unlikely that the millimeter emission is synchrotron radiation. Carilli et al. (\\cite{car01a}, \\cite{car01b}) show that for most of the high-redshift QSOs detected at 1.2 mm warm dust is the likely source of the emission.  \\subsection{FIR luminosity, dust mass, and star formation rate}  We therefore conclude that the millimeter flux of RD0301 is thermal in nature. When matching the observed 1.2~mm flux of 0.87 mJy to a grey body with a dust temperature of 50~K and an emissivity index $\\beta=1.5$ (well adapted for high-redshift sources -- see, e.g., Benford et al. 1999), the FIR luminosity of RD0301 is $L_{\\rm FIR}=4.0\\times 10^{12} \\, \\rm L_\\odot$, and the implied dust mass $\\approx 10^8 \\, \\rm M_\\odot$. The FIR luminosity of RD0301 is comparable to its optical luminosity, $L_{\\rm opt}\\approx 3.7\\times 10^{12}\\, \\rm L_\\odot$, given by the rest frame blue magnitude $M_{\\rm B} \\approx -24.0$ (for a $\\Lambda$-cosmology). We here adopted the blue luminosity bolometric correction factor of $\\sim 12$ derived for the PG sample by Elvis et al. (\\cite{elv94}).  Both the FIR luminosity and the dust mass are comparable to those derived for local infrared-luminous galaxies such as Arp~220 and Mrk~231 (e.g., Radford et al. \\cite{rad91}).  The thermal luminosity of RD0301 is comparable also to that of the bright high-redshift starburst galaxies found in deep optical surveys (Adelberger \\& Steidel \\cite{ade00}), and of those thought to be producing the bulk of the FIR background emission (Peacock et al. \\cite{pea00}).  If the FIR luminosity of RD0301 arose from a continuous starburst of age 10 to 100~Myr with a modified Salpeter IMF (see Omont et al.  \\cite{omo01} for details on the adopted starburst model), the implied star formation rate would be $\\approx 600 \\, \\rm M_\\odot~yr^{-1}$.  This value again is similar to those inferred for Arp~200 and Mrk~231.   \\subsection{Comparison with the SDSS and DPSS QSOs}  The millimeter flux of RD0301 is typical for the optically brightest high-redshift QSOs.  When averaging (as if we observed one single object) the MAMBO observations of the 112 redshift 3.6 to 5.0 QSOs observed in the Carilli et al. (\\cite{car01a}) SDSS and Omont et al. (\\cite{omo01}) DPSS surveys, one obtains an average flux density of $1.3\\pm 0.05$~mJy, which implies an average luminosity of $L_{\\rm FIR} \\approx 5\\times 10^{12} \\, \\rm L_\\odot$.  The dust luminosity of RD0301 is therefore close to the average dust luminosity of a sample of quasars with $-26>M_{\\rm B}>-29.5$, sources which are optically 10 to 100 times brighter than RD0301.  \\begin{figure}[] \\psfig{figure={fig2.eps},width=9cm} \\caption{MAMBO 1.2 mm flux density as a function of blue magnitude for the $z>3.6$ QSOs observed in the Carilli et al. (open symbols) and Omont et al. (filled symbols) surveys. Average source fluxes for the respective samples are indicated as large open squares, which are placed at the median $M_{\\rm B}$ of the respective sample.  } \\label{figure2} \\end{figure}  The mean blue magnitudes of the Carilli et al. (\\cite{car01a}) and Omont et al. (\\cite{omo01}) QSOs are $-27.0$ and $-27.9$ mag, respectively (Fig.~\\ref{figure2}).  Although the QSOs of these samples span a range of about 25 in optical brightness, the DPSS QSOs studied by Omont et al. (\\cite{omo01}) are on average twice as bright in the optical as the SDSS QSOs observed by Carilli et al. (\\cite{car01a}).  Similarly, the average 1.2~mm flux of the QSOs in the respective samples differs by a factor two, $0.92$ mJy versus $2.1$~mJy (large squares in Fig.~\\ref{figure2}). However, the significance of a possible correlation between the average optical and millimeter brightness of high $z$ QSOs remains low since the scatter in both quantities is much larger than the difference between the samples' averages (see Omont et al. [\\cite{omo02}] for a detailed statistical discussion).  RD0301 does not follow the possible trend of increasing millimeter flux with optical luminosity seen in the QSO millimeter surveys, as Fig. \\ref{figure2} well illustrates.  Further (sub)millimeter observations of optically faint ($M_{\\rm B}>-26$) high-redshift QSOs are needed to clarify the relation between their thermal and optical luminosities.  A lack of such a correlation would support the idea that the dust heating is caused by young, massive stars, and not by the AGN.  This hypothesis is also supported by the detection of CO emission in a number of high-redshift QSOs (Guilloteau \\cite{gui01}; Cox et al. \\cite{cox02}), and by the lack of strong radio emission from most of them (Carilli et al. \\cite{car01b}).  Alternatively, unification models of AGN might not necessarily predict a strong correlation between the orientation-dependent optical emission and the isotropic FIR flux from the AGN. But still the AGNs' FIR flux should correlate well with its hard X-ray emission, which does not suffer much from extinction. Millimeter-studies of high-redshift, hard X-ray sources would be most useful to address this interesting issue (e.g. Page et al. \\cite{pag01}).  SIRTF observations of mid-IR PAH features should also be able to clarify the origin of the FIR emission: a strong correlation between PAH and FIR emission supports the starburst origin of the FIR luminosity (Genzel et al. \\cite{gen98})."
    },
    "0201/astro-ph0201106_arXiv.txt": {
        "abstract": "We present the results of a  theoretical investigation aimed at testing whether full amplitude, nonlinear, convective models account for the I-band light curves of  Bump Cepheids in the Large Magellanic Cloud (LMC). We selected two objects from the OGLE sample that show a well-defined bump along the decreasing (short-period) and the rising (long-period) branch respectively. We find that current models do reproduce the luminosity variation over the entire pulsation cycle if the adopted stellar mass is roughly 15\\% smaller than predicted by evolutionary models that neglect both mass loss and convective core overshooting. Moreover, we find that the fit to the light curve of the long-period Cepheid located close to the cool edge of the instability strip requires an increase in the mixing length from 1.5 to 1.8 Hp. This suggests an increase in the efficiency of the convective transport when moving toward cooler effective temperatures. Current pulsation calculations supply a LMC distance modulus ranging from 18.48 to 18.58 mag. ",
        "introduction": "Classical Cepheids are key objects for estimating stellar distances and for calibrating secondary distance indicators necessary to evaluate the Hubble constant (Ferrarese et al. 2000; Saha et al. 1999). From a theoretical point of view Cepheids can also be regarded as fundamental physical laboratories to assess both the accuracy of the input physics (Simon 1989) and the reliability of the physical assumptions adopted in modeling radial pulsations (Bono, Marconi, \\& Stellingwerf 1999, hereinafter BMS). Thus, any new test concerning the accuracy of theoretical predictions appears of paramount relevance.  Moreover, this effort is mandatory to improve our confidence on Cepheids as a robust rung of the cosmic distance ladder and  to disentangle relevant open questions such as the dependence of the  Cepheid distance scale either on metal content (Bono et al. 1999; Bono, Castellani \\& Marconi 2000a; Caputo et al. 2000; Laney 2000; Storm et al. 2000) or on blending (Mochejska et al. 2000).  In this context, one has to remind that the current theoretical framework based on nonlinear, convective pulsation models, is facing with a fundamental problem, i.e. the calibration of the turbulent convection (TC) model adopted to account for the coupling between pulsation and convection (Castor 1968; Stellingwerf 1982). The approach originally suggested by Bono \\& Stellingwerf (1993, 1994, hereinafter BS) relies on the combination of leading physical arguments and empirical constraints on the topology of the RR Lyrae instability strip in Galactic globulars. Such a theoretical scenario accounts for some relevant properties of radial variables in the Cepheid instability strip (Bono et al. 1997a,b; BMS; Bono, Caputo, \\& Marconi 2001). However, the predicted RR Lyrae light curves close to the low-temperature instability edge are somehow at variance with empirical data (Kovacs \\& Kanbur 1998). This appears as a disturbing evidence, since empirical light curves supply tighter constraints on the accuracy of theoretical predictions than mean magnitudes and/or pulsational periods.  In a recent paper (Bono, Castellani,\\& Marconi 2000b, hereinafter BCM) we have approached such a problem and successfully reproduced the luminosity variation over a full pulsation cycle of the field, first overtone RR Lyrae \\uc. At the same time, we also constrained the adopted TC  model, and indeed in that paper we found that the Bump located just before the luminosity maximum can be reproduced by nonlinear models only by assuming a vanishing overshooting efficiency at the boundaries of the convective unstable regions. To explore the general reliability of this finding, we undertook a further investigation aimed at checking whether this TC model can also account for Cepheid light curves. To this purpose we selected two Bump Cepheids from the OGLE database (Udalski et al. 1999) of the LMC, with the Bump either along the decreasing (OGLE194103, short-period) or along the rising (OGLE56087, long-period) branch of the light curve. Table 1 lists the empirical data of the selected objects.  The reason for this choice follows from the evidence that both the amplitude and the phase of the Bump are crucial features to nail down the predicting power of theoretical models. Moreover, both stars present quite similar apparent magnitudes, and therefore the period difference suggests that the short-period Cepheid is located close to the blue (hot) edge, while the long-period one close to the red (cool) edge of the instability strip. When moving from hotter to cooler effective temperatures the thickness of the convective unstable region, and in turn the efficiency of the convective transport increases. This means that the envelopes of the two Bump Cepheids present substantially different physical structures. Finally, the intimate nature of the Hertzsprung progression is still controversial (Bono, Marconi, \\& Stellingwerf 2000c), and a reliable fit to empirical light curves can supply new insights into the physical mechanisms that govern the occurrence of this phenomenon. ",
        "conclusions": "According to the results given in the previous section, our pulsational computations require a larger mixing length value when the star becomes cooler.  We note that this is not a very surprising result, since a mixing length roughly equal to 1.9 Hp is also required by the evolutionary models that best fit the red giant branch of Galactic globular clusters (Castellani 1999). This notwithstanding, as already discussed in BCM, a substantial variation in the mixing length parameter causes a sizable change in the efficiency of the convective transport across the driving regions, and in turn in the temperature width of the instability strip. To estimate the dependence of the instability edges on {\\em l} we constructed new models with the stellar mass of the long-period best fit model. We find that when moving from {\\em l}=1.5 to 1.8 Hp the temperature width decreases from 1100 K to 800 K. This change is mainly due to a shift toward hotter effective temperatures of the red boundary and could be marginally at odds with empirical estimates. In fact, estimates by Pel \\& Lub (1978, see their Fig. 3) do suggest that in this period range the strip is slightly larger than 800 K. Unfortunately, we are not aware of any recent empirical estimates that can allow us to supply firm constraints on this observable.  The pulsational scenario we are dealing with accounts for the empirical light curves of the selected Bump Cepheids provided that their stellar masses are 15\\% smaller than predicted by canonical evolutionary models. This theoretical evidence, if taken at face value, could be due either to the occurrence of mass-loss among intermediate-mass stars, as originally suggested by Cox (1980), or to convective core overshooting, or both of them. Unfortunately, the pulsational approach adopted in this investigation does not allow us to discriminate between the two different hypothesis. Current results supply some support to the evidence recently brought out by Beaulieu et al. (2001) concerning the {\\em discrepancy} between evolutionary and pulsational masses among Magellanic Cepheids. On the basis of a new approach  they found that evolutionary masses are up to $\\Delta \\log M/M_\\odot = 0.1$ larger than pulsational ones. On the other hand, we found that for the two selected Bump Cepheids the discrepancy is of the order of 0.07 and 0.08 dex for the short and the long-period respectively.  Note that Bono et al. (2001) performed a detailed analysis of both pulsation and evolutionary masses of Galactic Cepheids for which are available both accurate empirical estimates of mean radii, distances, and individual reddenings. Interestingly enough, they found that pulsational masses are approximately 10\\% smaller than evolutionary ones. A similar result was found by Wood, Arnold, \\& Sebo (1997) by fitting the light curve of a LMC Bump Cepheid (HV905). However, they found that the luminosity of this object was substantially higher than predicted by ML relations based either on canonical or on mild overshooting evolutionary models.  To assess whether our best fit models support either the occurrence of a resonance between fundamental and second overtone or the Christy's echo mechanism (Bono et al. 2000c, and references therein) we finally constructed two new models by adopting the same input parameters of the models plotted in Fig. 1, but by perturbing the linear second overtone radial eigenfunctions. We find that both of them, after a transient phase lasting $\\approx$200 (long-period) and $\\approx$700 (short-period) cycles, undergo a mode switch from the second to the fundamental mode. The period ratios between second overtone and fundamental period are $P_2/P_0 (OGLE194103)$=0.523-0.528 and $P_2/P_0 (OGLE56087)$=0.495-0.503 respectively. For each object the former value was derived by adopting linear $P_2$ and $P_0$ periods, while the latter one with linear $P_2$ and nonlinear $P_0$ periods. The values of these ratios still fall in the so-called {\\em resonance region} (Simon \\& Schmidt 1976). However, we  note that the physical structure of the best fit models does not show any nodal line (Bono et al. 1997c) along the pulsation cycle. This finding together with the good agreement between theoretical predictions and empirical data further strengthens the plausibility that the Bumps are triggered by the Christy's echo mechanism."
    },
    "0201/astro-ph0201276_arXiv.txt": {
        "abstract": "{We present high-resolution rotation curves of a sample of 26 low surface brightness galaxies. From these curves we derive mass distributions using a variety of assumptions for the stellar mass-to-light ratio.  We show that the predictions of current Cold Dark Matter models for the density profiles of dark matter halos are inconsistent with the observed curves. The latter indicate a core-dominated structure, rather than the theoretically preferred cuspy structure.}  \\authorrunning{de Blok \\& Bosma} \\titlerunning{LSB galaxy rotation curves} ",
        "introduction": "\\label{sec:intro} \\indent  Despite much effort, it is still unclear to what extent rotation curves can give clues about the distribution of the visible and dark matter in bright spiral galaxies \\citep[e.g.][]{B99,Sellwood}.  It is thought, however, that low surface brightness (LSB) galaxies and dwarf galaxies are dark matter dominated, and that therefore the analysis of their rotation curves can yield directly important information about the properties and distribution of their associated dark matter halos \\citep{Blok97,Verheijen,Sw1999}.  This implication has far-reaching consequences, as early results by e.g. \\citet{Moore} already showed.  Predictions from cold dark matter (CDM) simulations were found to disagree with observations of rotation curves of several dwarf galaxies; the data indicating much less cuspy distributions of matter than the simulations. It was thought at the time that this problem could be solved once the effect of feedback due to star formation was understood. However, the behaviour of the rotation curves of low surface brightness (LSB) galaxies is rather similar, and their low star formation rate at present and in the past indicate that star formation in such galaxies might never have been important enough to modify their structure drastically.  Early work on LSB galaxies used 21-cm \\HI line work to determine the rotation curves \\citep[e.g.][]{Hulst,Blok96}, and as such the results are likely to suffer from modest angular resolution effects (commonly called beam smearing).  Even though these can be partly modelled, experience shows that direct measurements are preferable (cf.\\ discussions in \\citealt{B78,KGB}); in particular, supplementary data in the optical emission lines, such as H$\\alpha$ and [N{\\sc ii}] are always useful \\citep[cf.][]{vdK+B, Rub89, Corradi, Sw2000}. Thus a whole industry has sprung up to combine optical and 21-cm line rotation curves of all sorts of gas rich galaxies.  For the particular problem of the dark matter distribution in LSB galaxies \\citet*{Sw2000} presented supplementary H$\\alpha$ data for five LSB galaxies already observed in \\HI by \\citet*{Blok96}, and concluded that the influence of beam smearing on the \\HI curves was severe enough to question earlier conclusions regarding dark matter content and rotation curve shape of LSB galaxies. \\citet{MRdB} and  \\citet*{dBMR} (hereafter dBMR) reanalysed these data, and concluded that the discrepancy between H$\\alpha$ and \\HI data is only really significant for one of these five galaxies. Conclusions regarding the dark matter content and the shape of the dark matter distribution in LSB galaxies are thus not affected.  In a comparison of pseudo-isothermal and CDM halo models using high-resolution rotation curves of a sample of a further 29 LSB galaxies dBMR show that the so-called ``universal'' CDM halo-profile as parameterised in \\citet*{NFW-96} is not a good description of the data: the rotation curves generally show linear solid-body rise in the inner parts, which is inconsistent with the steeper rise necessary for the CDM rotation curves. Rather, the rotation curves prefer a pseudo-isothermal (i.e.\\ core-dominated) halo model. \\citet*{dBMBR} furthermore showed that \\emph{all} rotation curves of LSB galaxies measured so far are consistent with a pseudo-isothermal core model. Analyses that show that these rotation curves confirm the CDM NFW halo model \\citep[e.g.][]{BS2001} can be traced back to the fact that at lower resolutions the NFW and pseudo-isothermal models look sufficiently similar and the errors are large enough that a NFW model can be made to fit the data.  The higher-resolution data presented in \\citet*{dBMBR} now seem to have settled the observational side of the debate in favour of core-dominated LSB galaxy halos.  These results are independent of the stellar mass-to-light ratios one assumes in constructing these mass models.  It would therefore seem that current models of structure formation and galaxy evolution need to take into account the fact that most late-type galaxies have a constant density dark matter \\emph{core} rather than a cusp (see also results by \\citealt{Bor,Sal1,Sal2} for HSB galaxies).  As the signature of the core is clearest at small radii it is important to find galaxies where these inner parts are well-resolved and well-sampled.  In this paper we thus supplement the collection of high-resolution rotation curves of LSB galaxies from \\citet*{dBMR} with curves for an additional 26 galaxies, of which 12 are entirely new and 14 were already used in the analysis by \\citet*{dBMBR}. The 12 new galaxies in our sample have been specifically chosen to have small distances so that we can easily verify the discrepancies between the NFW model and the observations.  This paper is organised as follows: we describe our data in Section 2, and present the rotation curves in Section 3. Section 4 describes individual galaxies. In Section 5 we describe the derivation of the final rotation curves.  The mass models are then presented in Section 6. In Section 7 we discuss these models. Section 8 digresses into the consequences that systematic observational effects may or may not have on the data. In Section 9 we discuss the mass densities profiles inferred from the rotation curves. Section 10 summarises the paper. ",
        "conclusions": "We have presented high-resolution $H\\alpha$/\\HI rotation curves of a sample of 26 LSB and dwarf galaxies. We have fitted mass-models to these rotation curves assuming both a pseudo-isothermal (core-dominated) halo and a CDM NFW (cusp-dominated) halo. We find that the pseudo-isothermal halos generally provide better fits, though the difference is maybe not as pronounced as in dBMR, which can be traced back to the fact that our sample contains more galaxies of higher surface brightness than the dBMR sample.  We find more galaxies with low concentration parameters than predicted by numerical CDM simulations. An analysis of the mass-density profiles of the halos, as derived from minimum-disk rotation curves, shows that the galaxies in our sample are dominated by more-or-less constant density cores. As shown in dBMBR the few galaxies that show slopes consistent with CDM are usually the ones that are not well-resolved, so that one traces the edges of the constant-density cores rather than the cores themselves. We have illustrated this explicitly by comparing the mass-density profiles of a number of galaxies for which we have both high-resolution $H\\alpha$ curves and medium-resolution \\HI curves. In general the slopes derived from the high-resolution curves are less steep than those from the \\HI curves.  In conclusion, our high resolution rotation data on nearby dwarfs and LSB galaxies show that the halos of late-type LSB and dwarf galaxies are dominated by kpc-sized constant density cores inconsistent with the predictions of cuspy dark matter halos in cosmological numerical simulations."
    },
    "0201/astro-ph0201089_arXiv.txt": {
        "abstract": "Detailed examination of the balance between star-formation and nuclear activity in AGN and starburst galaxies leads to the composition of an Hertzsprung-Russell diagram in which possible evolutionary tracks can be drawn. It is likely that these tracks also relate to the level of obscuration. ",
        "introduction": "The question I'd like to address is whether there is a physical reason behind the resemblance of the cosmic star-formation history diagram, the `Madau-plot', and the evolving QSO space density diagram (e.g., Shaver et al.  1996).  Given that the shapes of these plots -- in particular the $z \\sim 2.5$ peaks -- are similar, could it be that star-formation and nuclear activity in galaxies are more intimately related than we have believed so far? Rosa Gonzalez Delgado and George Rieke (these Proceedings) have already presented evidence that such is the case for Seyfert galaxies -- I will add QSOs and discuss the AGN--starburst symbiosis within the framework of an Hertzsprung-Russell diagram analogon.  A more complete account of the issue and a more thorough presentation of the rationale can be found in Barthel (2001). ",
        "conclusions": "The 60\\,$\\mu$m luminosity, when normalized with the 25\\,$\\mu$m (or the 12\\,$\\mu$m), the blue optical or the radio luminosity, permits assessment of the absolute and relative strength of star-formation and nuclear activity in active galaxies and quasars.  The FIR temperature can be combined with measures of the bolometric luminosity to yield an intriguing AGN Hertzsprung-Russell diagram, which among other things suggests that star-formation plays an important role in many AGN.  Such photometric ratios can be obtained in a straightforwardly manner for the faint distant objects to be measured in large quantities with upcoming space-infrared missions, such as SIRTF, ASTRO-F and Herschel."
    },
    "0201/astro-ph0201040_arXiv.txt": {
        "abstract": "We have measured broadband optical $BVR$ photometry of 24 Classical and Scattered Kuiper belt objects (KBOs), approximately doubling the published sample of colors for these classes of objects.  We find a statistically significant correlation between object color and inclination in the Classical Kuiper belt using our data.  The color and inclination correlation increases in significance after the inclusion of additional data points culled from all published works. Apparently, this color and inclination correlation has not been more widely reported because the Plutinos show no such correlation, and thus have been a major contaminant in previous samples.  The color and inclination correlation excludes simple origins of color diversity, such as the presence of a coloring agent without regard to dynamical effects.  Unfortunately, our current knowledge of the Kuiper belt precludes us from understanding whether the color and inclination trend is due to environmental factors, such as collisional resurfacing, or primordial population effects.  A perihelion and color correlation is also evident, although this appears to be a spurious correlation induced by sampling bias, as perihelion and inclination are correlated in the observed sample of KBOs. ",
        "introduction": "Beyond Neptune, the Solar System is populated by $\\sim 10^{5}$ Kuiper belt objects (KBOs, also trans-Neptunians) larger than 100 km in diameter, which represent about 0.1 Earth masses of material and are presumed to contain the least thermally processed material in the solar system (Trujillo, Jewitt \\& Luu 2001).  As of Oct 2001, $\\approx 480$ KBOs are known, less than 1\\% of the inferred population. Dynamically, the KBOs can be divided into at least three classes: (1) the Classical KBOs represent 2/3 of the observed sample and about 1/2 the inferred population; (2) the Resonant KBOs are comprised primarily of Plutinos ($a \\approx 39.4$ AU) and objects in 2:1 mean-motion resonance with Neptune ($a \\approx 47.7$ AU), representing 15\\% of the observed sample, but only 6\\% of the total population; and (3) the Scattered KBOs comprise about 10\\% of the observed sample, but nearly half the inferred population because their extreme dynamical properties ($50 \\mbox{ AU} < a < 200 \\mbox { AU}$ and eccentricities $0.2 < e < 0.9$) place them at unobservably great heliocentric distances for large fractions of their orbits (Trujillo, Jewitt \\& Luu 2000).  Since the KBOs are thought to be the least thermally modified bodies in the solar system, they are expected to provide insight into the physical and chemical processes that dominated the primitive solar nebula.  Given the fact that the KBOs have been observed for nearly a decade, very little is known about their physical properties, primarily because of their extreme faintness (median $R$ magnitude $m_{R} = 23.0$).  The most striking physical feature observed of the KBOs is their substantial color diversity, from nearly neutral reflectance to the reddest objects in the solar system (Jewitt \\& Luu 1998 and Tegler \\& Romanishin 2000).  Objects are reddened from the visible throughout the near-IR (Davies et al. 2000), presumably from a single coloring agent (Jewitt and Luu 2001), although the physical and chemical mechanisms for the reddening process have not been identified.  Recently, Tegler \\& Romanishin (2000) reported that the Classical KBOs exhibited a correlation between color and inclination. No other observers have reported this trend, and Tegler \\& Romanishin estimate the significance of their reported trend using their two-color paradigm, which has not been confirmed by independent sources (Barucci et al. 2001 and Jewitt \\& Luu 2001).  Thus, we chose to conduct our own $BVR$ observations of this population.  Using our new observations, we demonstrate that the inclination and color correlation is indeed statistically significant.  In addition, we note the presence of a perihelion and color correlation which appears to be due to sampling bias. ",
        "conclusions": ""
    },
    "0201/astro-ph0201510_arXiv.txt": {
        "abstract": "We present results from a deep mid--infrared survey of the Hubble Deep Field South (HDF--S) region performed at 7 and 15$\\mu$m with the CAM instrument on board the {\\em Infrared Space Observatory} ({\\em ISO}). We found reliable optical/near--IR associations for 32 of the 35 sources detected in this field by Oliver et al. (2002, Paper I): eight of them were identified as stars, one is definitely  an AGN, a second seems likely to be an AGN, too, while the remaining 22 appear to be normal spiral or starburst galaxies. Using model spectral energy distributions (SEDs) of similar galaxies, we compare methods for estimating the star formation rates (SFRs) in these objects, finding that an estimator based on integrated (3--1000$\\mu$m) infrared luminosity reproduces the model SFRs best. Applying this estimator to model fits to the SEDs of our 22 spiral and starburst galaxies, we find that they are forming stars at rates of  $\\sim1 - 100$ $M_{\\odot} \\; {\\rm yr}^{-1}$, with a median value of $\\sim 40$ $M_{\\odot} \\; {\\rm yr}^{-1}$, assuming an Einstein -- de Sitter universe with a Hubble constant of 50~km~$s^{-1}$~Mpc$^{-1}$, and star formation taking place according to a Salpeter (1955) IMF across the mass range $0.1-100M_{\\odot}$. We split the redshift range $0.0 \\leq z \\leq 0.6$ into two equal--volume bins to compute raw estimates of the star formation rate density,  $\\dot{\\rho}_*$,  contributed by these sources, assuming the same cosmology and IMF as above and computing errors based on estimated uncertainties in the SFRs of individual galaxies. We compare these results with other estimates of $\\dot{\\rho}_*$ made with the same assumptions, showing them to be consistent with the results of Flores et al. (1999) from their {\\em ISO\\/} survey of the CFRS 1415+52 field. However, the relatively small volume of our survey means that our $\\dot{\\rho}_*$ estimates suffer from a large sampling variance, implying that our results, by themselves, do not place tight constraints on the global mean star formation rate density. ",
        "introduction": "In an accompanying paper (Oliver et al. 2002, Paper I) we described the survey we made of the Hubble Deep Field South (HDF--S\\footnote{see {\\tt www.stsci.edu/ftp/science/hdfsouth/hdfs.html}}) region using the LW2 (centred at 6.7$\\mu$m) and LW3 (15$\\mu$m) filters of the CAM (Cesarsky et al. 1996) instrument on board the {\\em Infrared Space Observatory (ISO)\\/} (Kessler et al. 1996). A prime motivation for this project came from the results of previous {\\em ISO\\/} surveys, such as our own survey (Serjeant et al. 1997, Goldschmidt et al. 1997, Oliver et al. 1997, Mann et al. 1997, Rowan--Robinson et al. 1997) of the northern Hubble Deep Field (HDF--N, Williams et al. 1996) and that of Flores et al. (1999) in the Canada--France Redshift Survey 1415+52 field (Lilly et al. 1995), which revealed an infrared luminosity density at $0.5 \\la z \\la 1$ implying a higher star formation rate density at those redshifts than indicated by previous optical studies (e.g. Lilly et al. 1996, Madau et al. 1996, Connolly et al. 1997).  These studies were facilitated by the availability of multi--wavelength datasets in those well--studied fields, which made possible the association of {\\em ISO\\/} sources with galaxies for which redshifts had been determined through optical spectroscopy or for which they could be estimated on the basis of multi--band optical/near--infrared photometry. The same factors favour the study of the HDF--S region, and in this paper we discuss the association of our {\\em ISO\\/} sources in the HDF--S with objects in optical and near--infrared catalogues of that field, and the estimation of their star formation rates. This process is significantly easier than was the case for our survey of the northern HDF--N, since, as emphasised in Paper I, our {\\em ISO\\/} dataset in the HDF--S is greatly superior, as a result of an observational strategy and data reduction procedure both revised significantly in the light of developing knowledge of the characteristics of the CAM instrument and of the source population it probed.  The plan of the rest of this paper is as follows. In Section~\\ref{sect:associations}, we describe the existing optical and near--infrared catalogues in the HDF--S region with which we shall seek associations for our {\\em ISO\\/} sources, review the likelihood ratio method used to obtain them and present the results of its application. Section~\\ref{sect:sfr_methods} compares a variety of commonly used star formation rate (SFR) estimators, through applying them to model SEDs, while, in  Section~\\ref{sect:sed_sfr}, we apply the best of them to the multi--wavelength photometric datasets compiled for the galaxies with which we have associated our {\\em ISO\\/} sources, to yield estimates of their SFRs, and Section~\\ref{sect:madau} uses these results to assess the contribution of {\\em ISO\\/}--selected sources to the star formation history of the HDF--S. Finally, in Section~\\ref{sect:discuss}, we discuss the results of this paper and the conclusions we draw from them. Appendices present mathematical details of the processes of correcting an SFR estimate to a canonical IMF, and of computing the sampling variance in estimates of the SFR density using a lognormal model for the cosmological density field.  \\begin{figure} \\epsfig{file=mb512rv_fig1.ps,angle=0,width=10cm} \\caption{This figure shows the location of our ISO rasters with respect to those of other datasets taken in the area. The shaded regions mark the {\\em HST} fields, with the STIS and NICMOS fields to the east and south of the WFPC2 field, respectively. The thick--dashed irregular shape and the solid circle show, respectively, the maximum extent of our ISO coverage (the coverage in the two bands differs slightly) and the region from which the source catalogues of Paper I were selected. The remaining lines show boundaries of four optical/near-IR surveys discussed in Section \\protect\\ref{sect:associations}, as follows: (i) dotted line -- AAT prime focus imaging survey of Verma et al.  (2002); (ii) dashed line -- CTIO BTC survey of Gardner et al. (1999); (iii) dot-dashed line -- CTIO BTC survey of \\protect\\cite{Walker 1999}; (iv) dot-dot-dot-dashed line -- ESO EIS optical imaging survey of Da Costa et al. (1998); and (v) long dashed line -- ESO EIS near-infrared survey of da Costa et al. (1998).} \\label{fig:fields} \\end{figure} ",
        "conclusions": "\\label{sect:discuss}  In this paper and its companion (Paper I) we have presented results from our {\\em ISO\\/} survey of the Hubble Deep Field South. Here we sought optical identifications for the {\\em ISO\\/} sources found in Paper I, obtaining reliable associations for 32 out of the 35 sources: these associations should be more secure than those made by Mann et al. (1997) using our initial analysis of our corresponding {\\em ISO\\/} survey of the northern Hubble Deep Field, thanks to the inclusion of corrections for CAM image distortions, which were not well characterised in 1997. Of these 32 sources, a total of twenty two were identified as spiral or starburst galaxies, eight of which have spectroscopic redshifts (from the work of Rigpoulou et al. 2000 and Glazebrook et al. 2002) and the remaining fourteen of which have had photometric redshifts estimated by ourselves and others. We estimate that our photometric redshifts should be accurate to $\\delta z =0.1$ or so, on the basis of the set of eight objects for which spectroscopic redshifts are known, although we have noted individual cases where the error is likely to be greater than that, for example when a redshift has been estimated solely on the basis of JHK photometry. We found that the redshift distribution of the galaxies associated with our {\\em ISO\\/} sources is consistent with that for similarly--bright optical galaxies in the HDF--S region as a whole.  We reviewed a series of methods commonly used to determine star formation rates for galaxies in  survey data, typically comprising a single broad--band flux for each object. We assessed the ability of these to reproduce the SFRs of models of actively star--forming galaxies, finding, as others (e.g. Granato et al. 2000) have, that those probing their far--infrared emission fared much better than those based on UV luminosity, indicating that massive star formation in these galaxies takes place in dusty regions: the work of Rigopoulou et al. (2000) shows further that this dust obscuration is not well corrected for using optically--determined extinction values. All these methods look for indications of the formation of massive stars, so a further complication in the determination of the absolute rate at which stellar mass is being formed in a given galaxy is uncertainty in the IMF, since the vast majority of the mass resides in stars not directly probed by these methods, so SFR estimates differing by factors of a few can result if differing IMFs are assumed. Caution must be exercised when applying these estimators in situations where the SED type of the galaxies under study are constrained only by a small number of broad--band fluxes, suggesting that constraints on the cosmic star formation history are best obtained in well--studied fields, such as the Hubble Deep Fields, where rich, multiwavength datasets are available. The small areas of the Hubble Deep Fields do, however, mean that significant sampling variances exist for estimates of $\\dot{\\rho}_*$ at $z \\leq 1$, and we showed, via two simple methods for assessing their magnitude, that while our results are consistent with those of previous authors, notably Flores et al. (1999) and Haarsma et al. (2000), they cannot by themselves yield tight constraints on the star formation history of the Universe, due to these sampling effects. Further details of this project can be found at {\\tt astro.ic.ac.uk/hdfs}."
    },
    "0201/astro-ph0201383_arXiv.txt": {
        "abstract": "We present spectral fits and timing analysis of Rossi X-ray Timing Explorer observations of \\gx. These observations were carried out over a span of more than two years and encompassed both the soft/high and hard/low states. Two observations were simultaneous with Advanced Satellite for Cosmology and Astrophysics observations.  Hysteresis in the soft state/hard state transition is observed.  The hard state exhibits a possible anti-correlation between coronal compactness (i.e., spectral hardness) and the covering fraction of cold, reflecting material.  The correlation between `reflection fraction' and soft X-ray flux, however, appears to be more universal.  Furthermore, low flux, hard state observations--- taken over a decline into quiescence--- show that the Fe line, independent of `reflection fraction', remains broad and at a roughly constant equivalent width, counter to expectations from Advection Dominated Accretion Flow models.  All power spectral densities (PSD) of the hard state X-ray lightcurves are describable as the sum of just a few broad, quasi-periodic features with frequencies that roughly scale as coronal compactness, $\\ell_{\\rm c}$, to the $-3/2$ power.  This is interpretable in a simple, toy model of an \\emph{efficient} spherical corona as variations of $\\ell_{\\rm c} \\propto R_{\\rm t}$, where $R_{\\rm t}$ is the `transition radius' between the corona and an outer thin disc.  Similar to observations of Cyg X-1, time lags between soft and hard variability anti-correlate with coronal compactness, and peak shortly after the transition from the soft to hard state.  A stronger correlation is seen between the time lags and the `reflection fraction'. These latter facts might suggest that the time lags are associated with the known, spatially very extended, synchrotron emitting outflow. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0201/astro-ph0201456_arXiv.txt": {
        "abstract": "Europium isotopic abundance fractions are reported for the very metal-poor, neutron-capture-rich giant stars CS~22892-052, HD~115444, and \\bd17. The abundance fractions, derived from analysis of several strong \\ion{Eu}{2} lines appearing in high-resolution spectra of these stars, are in excellent agreement with each other and with their values in the Solar System: fr(\\iso{Eu}{151})~$\\simeq$ fr(\\iso{Eu}{153})~$\\simeq$ 0.5. Detailed abundance studies of very metal-poor stars have previously shown that the total elemental abundances of stable atoms with atomic numbers z~$\\geq$~56 typically match very closely those of a scaled solar-system $r$-process abundance distribution. The present results for the first time extend this agreement to the isotopic level. ",
        "introduction": "Neutron bombardment reactions create all abundant isotopes of more than 60 elements of the periodic table. Neutron-capture ($n$-capture) nucleosynthesis can occur slowly in the $s$-process, in which $\\beta$ decays have time to happen between successive neutron captures, or rapidly in the $r$-process, in which enormous, but short-lived, neutron fluxes temporarily overwhelm $\\beta$-decay rates. After completion of an $n$-capture event, the final isotopic abundance mix will be very different for the $r$- and $s$-processes. Detailed exploration of these $n$-capture events can best be done with full knowledge of the isotopic abundances in stars over a large metallicity range. At present, this can be done for one star, the Sun, through the isotopic analysis of carbonaceous chondrite meteorites. Such analysis reveals a complex mix of $r$- and $s$-process contributions to Solar System $n$-capture elements.  Meteoritic data can be obtained only for the Solar System, so $n$-capture contents of other stars must be gleaned from their spectra. However, the isotopic wavelength shifts for transitions of almost all $n$-capture elements are small compared to the thermal+turbulent line widths in stellar spectra. Thus isotopic abundance information is usually inaccessible, and confrontation of $n$-capture predictions can only be done with total elemental abundances. In the most metal-poor stars, these abundances are very different than their total ($s$- plus $r$-process) Solar System values.  Neglecting the cases of anomalously carbon-rich stars, abundances of $n$-capture elements with Z~$\\geq$ 56 in extremely metal-poor stars are consistent only with the $r$-process component abundances in Solar-System material (e.g. McWilliam 1997\\nocite{McW97}; Burris et~al. 2000\\nocite{BPASCR00}; Sneden et~al. 2000\\nocite{SCIFBBL00}; Westin et~al 2000\\nocite{WSGC00}; Cayrel et~al. 2001\\nocite{Cay01}). This agreement suggests that even though the $r$-process is a very violent dynamical event that cannot be reproduced in the laboratory, and is very difficult to model theoretically, nature may be constrained to produce $r$-process abundances in basically one fashion. But the case could be strengthened if the elemental abundances could be deconvolved into isotopic fractions. In this {\\em Letter} we report a step in this direction, by demonstrating that abundance fractions of europium's two stable isotopes are essentially equal in three $n$-capture-rich halo stars, just as they are in the Solar System. ",
        "conclusions": "Both europium isotopes, \\iso{Eu}{151}\\ and \\iso{Eu}{153}. are produced primarily by the {\\it r}-process in Solar System material (see K{\\\"a}ppeler et al. 1989), and thus this element has only a very small {\\it s}-process contribution (Burris et~al. 2000\\nocite{BPASCR00}). These conclusions are based upon employing the {\\it classical} or {\\it standard} model for predicting {\\it s}-process production and the associated {\\it r}-process Solar System residuals. More complex models have been developed to predict {\\it s}-process abundances in low- to intermediate-mass asymptotic giant branch (AGB) stars, the likely dominant site for this type of synthesis. However, even in these stellar models the calculated {\\it s}-process abundances for \\iso{Eu}{151}\\ and \\iso{Eu}{153} are small, and the abundance ratio remains solar (see Arlandini et~al. 1999\\nocite{AKWGLBS99}).  We can also rule out large amounts of $s$-process contribution to the europium abundances of our program stars, either by past mass transfer from nearby AGB stars or from internal nucleosynthesis. Such events would yield large $s$-process contributions to a number of elements, not just to europium. Consider, for example, the neighboring elements barium and lanthanum, whose synthesis in Solar System matter has been dominated by the $s$-process. The mean observed abundance ratios of these elements with respect to europium for the program stars, from results of Sneden et~al. (2000), Westin et~al. (2000), and Cowan et~al. (2001), are $<$log~$\\epsilon$(Ba/Eu)$>$~$\\simeq$ +1.0 and $<$log~$\\epsilon$(La/Eu)$>$~$\\simeq$ +0.2. But in the Solar System these ($r$- plus $s$-process) abundance ratios are +1.8 and +0.7, respectively (Burris et~al. 2000\\nocite{BPASCR00}). Moreover, $s$-process computations by Malaney (1987\\nocite{Mal87}) covering a large range of neutron exposure parameters yield $minumum$ values for the ratios of +2.5 and +1.3, respectively. It cannot be simultaneously argued that the europium abundances in our stars have significant $s$-process components while their Ba/Eu and La/Eu elemental abundance ratios are more than an order of magnitude smaller than in the Solar System or in $s$-process predictions.\\footnote{ One abundance anomaly among the three target stars should be noted. The carbon abundance of CS~22892-052 is very large ([C/Fe]~$\\simeq$ +1.0: McWilliam et~al. 1995\\nocite{MPSS95}; Norris, Ryan, \\& Beers 1997\\nocite{NRB97}) compared to that of the other two stars ([C/Fe]~$<$ 0: Kraft et~al. 1982\\nocite{Ketal82}; Westin et~al. 2000; Cowan et~al. 2001). Carbon overabundances in some very metal-poor stars usually are accompanied by significant $s$-process abundance enhancements (e.g, Norris et~al. 1997\\nocite{NRB97}; Hill et~al. 2000\\nocite{Hetal97}). However, some carbon-rich low metallicity stars are known to be $s$-process-poor (Preston \\& Sneden 2001\\nocite{PS01}; Aoki et~al. 2001\\nocite{ANRBA01}). Probably the large carbon abundance and $n$-capture enhancements of CS~22892-052 were generated in different nucleosynthesis events.}  One further consideration would be whether the {\\it r}-process itself might not always produce Solar System isotopic abundance ratios for europium. Those abundances are determined, for example, by the neutron number density and the corresponding flow ``path'' for this synthesis. However. such a variation would also affect other isotopes, producing non-solar elemental abundance ratios -- something that is not observed in very metal-poor halo stars. Recent {\\it r}-process models (e.g., Pfeiffer, Kratz, \\& Thielemann 1997\\nocite{PKT97}; Cowan et~al. 1999\\nocite{CPKTSBTB99}) that reproduce the observed solar {\\it r}-process elemental abundances in these stars also replicate the Solar System isotopic ratios for elements such as europium.  The observed Solar System ratio of the europium isotopes in our sample of three low-metallicity, old halo stars suggests that successful {\\it r}-process models will need to be those that match the Solar System {\\it r}-process abundances, at least for the heaviest neutron-capture elements. Whether there is a narrow range of astrophysical and nuclear conditions (Cameron 2001\\nocite{Cam01}), a narrow mass range for the sites of the r-process (Mathews et~al. 1992\\nocite{MBC92}; Wheeler et~al. 1998\\nocite{WCH98}), or some other constraining factors, there appears to be a uniformity in the production of these elements and, at least for europium, the isotopes. Several groups have high resolution and S/N data for other very metal-poor stars with enhanced rare-earth elemental abundances. Further Eu isotopic abundance analyses should be undertaken to determine if there are exceptions to the apparent equality of Eu isotopic fractions in low metallicity halo stars."
    },
    "0201/hep-th0201127_arXiv.txt": {
        "abstract": "For the 2-brane Randall-Sundrum model, we calculate the bulk geometry for strong gravity, in the low matter density regime, for slowly varying matter sources. This is relevant for astrophysical or cosmological applications.  The warped compactification means the radion can not be written as a homogeneous mode in the orbifold coordinate, and we introduce it by extending the coordinate patch approach of the linear theory to the non-linear case.  The negative tension brane is taken to be in vacuum. For conformally invariant matter on the positive tension brane, we solve the bulk geometry as a derivative expansion, formally summing the `Kaluza-Klein' contributions to all orders.  For general matter we compute the Einstein equations to leading order, finding a scalar-tensor theory with $\\omega(\\Psi) \\propto \\Psi / (1 - \\Psi)$, and geometrically interpret the radion.  We comment that this radion scalar may become large in the context of strong gravity with low density matter. Equations of state allowing $(\\rho - 3 P)$ to be negative, can exhibit behavior where the matter decreases the distance between the 2 branes, which we illustrate numerically for static star solutions using an incompressible fluid.  For increasing stellar density, the branes become close before the upper mass limit, but after violation of the dominant energy condition.  This raises the interesting question of whether astrophysically reasonable matter, and initial data, could cause branes to collide at low energy, such as in dynamical collapse. ",
        "introduction": "Much progress has been made in understanding the long range gravitational response of branes at orbifold fixed planes, to localized matter. The Randall-Sundrum models, with one or two branes \\cite{Randall:1999ee,Randall:1999vf,Lykken:1999nb} simply have gravity and a cosmological constant in the bulk, making the linear problem very tractable. Linear calculations \\cite{Garriga:1999yh,Giddings:2000mu,Tanaka:2000zv,Sasaki:1999mi} show that in the one brane case, long range $4$-dimensional gravity is recovered on the brane. Little is known of the general non-linear behavior \\cite{Emparan:1999wa,Emparan:2001ce}.  There is no mass gap, and thus the non-linear problem is essentially a 5-dimensional one, even for long wavelength sources.  Second order perturbation theory \\cite{Kudoh:2001wb,Giannakis:2000zx}, and fully non-linear studies \\cite{Wiseman:2001,Gregory:2001xu} are again consistent with recovering usual 4-dimensional strong gravity.  In particular \\cite{Wiseman:2001} shows that non-perturbative phenomena, such as the upper mass limit for static stars, extend smoothly from large to small objects, whose characteristic sizes are taken relative to the AdS length.  For the two brane case, linear theory \\cite{Garriga:1999yh} found that the effective gravity is Brans-Dicke, for an observer on the positive tension brane, and a vacuum negative tension brane.  For reasonable brane separations phenomenologically acceptable gravity is recovered. For observers on the negative tension brane it was found the response was incompatible with observation for any brane separation, although stabilizing the distance between the branes does allow one to recover standard 4-dimensional gravity \\cite{Goldberger:1999uk,Goldberger:1999un,Goldberger:1999wh,Csaki:1999mp,Tanaka:2000er,Mukohyama:2001ks,Kudoh:2001kz}. In this paper we consider an observer on the positive tension brane, the negative tension brane to be in vacuum, and the orbifold radius to be unstabilized. This allows us to study the dynamics of the radion in strong gravity. Of course, our methods may be extended to include the stabilized case too. Although moduli have previously been assumed to be fixed at late times, a crucial feature of the recent cyclic Ekpyrotic scenario \\cite{Steinhardt:2001st,Steinhardt:2001vw}, is that the radius of the orbifold is not stabilized, the cyclic nature of this model being such that the separation always remains finite.  The two brane strong gravity case appears more tractable than for one brane, as the linear theory allows a Kaluza-Klein style reduction of the propagator \\cite{Garriga:1999yh}. In Kaluza-Klein compactifications, only the homogeneous zero modes are excited on long wavelengths, and the matter is thought to comprise fields over the whole internal space.  However, for matter to be supported on the orbifold branes there must be modes excited which are not homogeneous in the extra dimension. Thus one cannot simply write down a non-linear ansatz for the metric, as is familiar from Kaluza-Klein compactifications.  In previous work on the Horava-Witten compactification \\cite{Horava:1996ma}, the low energy, long range effective theory was constructed using consistent orbifold reductions \\cite{Lalak:1998ti,Lukas:1998fg,Lukas:1998tt,Lukas:1998ew,Lukas:1999yn}. The methods developed treated the homogeneous component of the metric as a background, and inhomogeneous perturbations about it as the contribution of the massive `Kaluza-Klein' modes. The metric was then solved to leading order in a derivative expansion. Strong gravity was not discussed in these works except when considering inflation \\cite{Lukas:1999yn}. In order to use these constructions, all the zero modes of the orbifold must be homogeneous. One important result of the warped compactification, is that whilst the graviton modes are homogeneous, the radion zero mode is not \\cite{Charmousis:1999rg}. This appears to be a general feature of warped models \\cite{Gen:2000nu,Chacko:2001em,Kogan:2001qx}, and one cannot directly use these Horava-Witten methods.  The aim of this paper is to apply the orbifold reduction to the warped Randall-Sundrum model, and then examine the non-linear behavior of the radion. We begin in section \\ref{sec:scalar} by illustrating the reduction proposed in \\cite{Lalak:1998ti,Lukas:1998fg,Lukas:1998tt,Lukas:1998ew,Lukas:1999yn} with a field theory example and discuss its application to strong gravity. In order to apply the method we must find a way to include the inhomogeneous radion zero mode. Instead of explicitly using an ansatz including the radion, we include it by `deflecting' the brane relative to a non-linear extension of the `Randall-Sundrum gauge' \\cite{Garriga:1999yh,Giddings:2000mu}. In sections \\ref{sec:metric} and \\ref{sec:conf} we solve the bulk metric for conformally invariant, low energy, long wavelength matter on the positive tension brane. For conformally invariant matter, this coordinate system is Gaussian normal to the brane. The bulk geometry is constructed using a derivative expansion, which is formally summed to all orders.  In section \\ref{sec:nonconf}, we extend the linear analysis to the non-linear, for general low density, long wavelength matter, allowing the brane to become deflected relative to this coordinate system. This scalar deflection becomes the radion, and is treated non-linearly.  This is crucial, as this field takes values roughly of order the Newtonian potential, for static systems.  Therefore, the case of strong gravity is exactly where the radion must be treated non-linearly.  We calculate a local action for the effective induced Einstein equations, to leading order in the derivative expansion.  The resulting theory is a scalar-tensor model, which reduces to Brans-Dicke in the linear approximation. This is the first non-linear, covariant derivation of the effective action.  Considering the zero modes \\cite{Chiba:2000rr} does not tell one the conformal metric that matter couples to. Whilst cosmological solutions such as \\cite{Csaki:1999mp,Goldberger:1999un,Binetruy:1999hy,Binetruy:2001tc} may derive radion dynamics, this only applies for a homogeneous radion, and thus is not a covariant derivation.  Indeed the examples \\cite{Csaki:1999mp,Goldberger:1999un} illustrate this point well.  The form of the metric chosen, whilst coincidentally reproducing the correct effective action \\cite{Chiba:2000rr}, does not solve the correct linearized equations \\cite{Charmousis:1999rg}, and thus only applies to the case of cosmological symmetry.  We discuss in detail the validity of the approximation for considering low density, long wavelength strong brane gravity. Terms that are quadratic in the energy density are neglected in the orbifold reduction method, consistent with the low energy approximation. The characteristic length scale of the matter must be large compared to the compactification scale for the derivative expansion to be valid. For strong gravity, the approximation then holds provided the 4-dimensional induced brane curvature invariants are small, compared to the compactification scales. Thus this method will not allow a global solution to a black hole geometry, or any other spacetimes with curvature singularities.  However, it will apply to all other cases of strong gravity, such as static relativistic stars, an example being neutron stars, or dynamical non-linear systems, such as binary neutron star systems, or collapse of matter up to the point where curvatures become singular.  Having pointed out that the radion may take large values for strong gravity configurations, we illustrate this using the example of relativistic stars.  Previous work \\cite{Germani:2001du,Deruelle:2001fb} has considered stars on branes using the projection formalism of \\cite{Shiromizu:1999wj}, allowing the quadratic stress tensor corrections to be calculated by making ansatzes on the bulk geometry. However, these corrections are assumed to be negligible in our low energy density assumption, and it is the bulk geometric corrections which are relevant, and cannot be calculated in the projection approach. In \\cite{Wiseman:2001} the full bulk solutions were numerically constructed for the one brane Randall-Sundrum case.  Using the methods here, we are able to analytically construct the geometry for large stars in the 2-brane case. In section \\ref{sec:star}, using the leading order action, we numerically consider the static relativistic star, for incompressible fluid matter.  The linear theory shows that for perturbative stars with positive density the branes are deflected apart.  However, it indicates that if $\\rho - 3 P$ becomes negative the opposite could occur. Now understanding the radion and bulk geometry non-linearly, we consider this, finding that as non-linear effects become important the branes do indeed become closer.  Furthermore, the branes appear to meet \\emph{before} the upper mass limit is reached. This will occur for any brane separation.  However, for phenomenologically acceptable separations \\cite{Garriga:1999yh}, we find it does not occur before the dominant energy condition is violated. Neutron stars are believed to have polytropic equations of state that do not support negative $\\rho - 3 P$. It then remains a very interesting, and tractable $1+1$ problem to understand whether dynamical systems may cause branes to collide, for realistic matter and initial data. We then have the possibility that physical matter at low energies and curvatures, compared to compactification scales, might cause brane collisions requiring a Planck energy physics description. ",
        "conclusions": "We have extended the linearized analysis of the Randall-Sundrum 2 brane compactification to allow the bulk, and induced geometry, to be calculated in the case of low energy strong gravity, with an unperturbed, vacuum negative tension brane. The induced geometry must have low curvature compared to the compactification curvature scale, and the matter energy density must be slowly varying on scales much larger than the compactification length. This is the situation relevant for non-linear gravitation in late time cosmology or astrophysics.  The radion is inhomogeneous in the transverse coordinate for warped compactifications. It therefore can not be included as a homogeneous zero mode, as in usual orbifold reductions. We show how to extend the coordinate systems used in the linear analysis, to the non-linear, low energy case. Firstly we choose a coordinate system analogous to the linear `Randall-Sundrum' gauge. This is Gaussian normal to branes with conformally invariant matter. We then solve the bulk geometry for conformally invariant matter on the positive tension brane, using a formal derivative expansion, which we sum to all orders. General matter is accommodated by introducing a non-linear `deflection' of the positive tension brane with respect to this coordinate system. This deflection is governed by the radion, and is determined to all orders in the bulk derivative expansion. We are then able to understand the radion geometrically, as the proper distance between the two orbifold branes, along a geodesic normal to the vacuum negative tension brane.  Whilst we have applied the method to the 2-brane Randall-Sundrum geometry, the methods can be extended to any warped compactification where the radion is not a simple homogeneous mode. Note also, that we have considered a vacuum negative tension brane, allowing only matter on the positive brane. This is purely for convenience, and can be relaxed more generally using the same techniques.  Having solved the bulk geometry, the effective action on the brane is calculated. To leading order, this yields a scalar-tensor theory with $\\omega \\propto \\Psi / (1 - \\Psi)$, where $\\Psi$ is the scalar depending on the distance between the branes. We have noted that the radion scalar may take values of order $\\sim O(1)$ in a strong gravity context, allowing the branes to touch or collide when gravitational fields become strong, and yet curvatures remain small compared to the compactification scale curvature. We have performed calculations for static spherical stars using incompressible fluid matter. For very low density stars, where the response remains linear, the branes are deflected apart. For densities where the gravitational response becomes strong, the brane separation may become reduced. We give evidence that the branes will touch before the stellar upper mass limit is reached, although the dominant energy condition is certainly violated. We find that, as indicated in the linear theory, $\\rho - 3 P$ must become negative to get an appreciable effect. Whilst this does occur for incompressible fluid, it is thought not to for physically reasonable stellar material.  Thus we find unrealistic matter is required to create touching branes in a static context. It raises the interesting question of whether realistic matter may allow branes to locally collide in a dynamic context. If this were the case, high energy physics may be required to understand low energy astrophysical processes."
    },
    "0201/hep-ph0201163_arXiv.txt": {
        "abstract": "The final stage of Bose-Einstein condensation in a large volume occurs through coarsening---growth of individual correlated patches. We present analytical arguments and numerical evidence that in the momentum space this growth corresponds to strong turbulence of the particle number, with a turbulent cascade towards the infrared and the power law $n(k) \\propto k^{-(d+1)}$ in $d=2,3$ spatial dimensions. As a corollary, we find that the correlation length grows linearly in time, with the speed of order of the speed of Bogoliubov's quasiparticles (phonons). We use these results to estimate the speed of BEC in atomic vapors and in galactic dark matter. \\\\ \\hspace*{3.9in} hep-ph/0201163 ",
        "introduction": "The discovery of Bose-Einstein condensation (BEC) in alkali vapors \\cite{exp1}--\\cite{exp3} had stimulated interest in a theoretical description of non-ideal Bose gases, a research topic that dates back to Bogoliubov's work \\cite{Bogoliubov} on the equilibrium properties. In current experiments, the gas is confined in a small magnetic trap and is strongly affected by the confining potential. However, theorists have also investigated the kinetics of BEC in a uniform (on average) gas, contained in a macroscopic volume \\cite{Kraichnan}--\\cite{grinst}. In addition to the intrinsic interest that this problem has for basic science, it is also of potential importance for astrophysics and cosmology in connection with the problem of galactic dark matter \\cite{it86}--\\cite{Riotto&Tkachev}.  The earliest theoretical work we are aware of on this subject is the paper by Kraichnan \\cite{Kraichnan}. Although we will confirm below several qualitative results of that work (such as concentration of particles in a progressively smaller momentum range and a turbulent cascade at larger momenta), the quantitative details differ. We attribute the difference to the imprecise treatment of phonons in ref. \\cite{Kraichnan} (noted by Kraichnan himself). Although the amount of theoretical work on BEC in a macroscopic volume has grown rapidly in recent years \\cite{KSS}--\\cite{grinst}, an adequate understanding of the phonon dynamics has remained elusive. In the present paper, we attempt to improve the situation, through a combination of analytical estimates and numerical results.  We consider an isolated nonrelativistic Bose gas, sufficiently dilute for the interaction between the particles (atoms in the laboratory, or dark matter particles in the outer space) to be characterized by a single parameter---the scattering length. Such a gas can be described by an effective Hamiltonian: \\be H = \\frac{\\hbar^2}{2m} \\nabla \\psi^{\\dagger} \\nabla\\psi + \\half g_4 (\\psi^{\\dagger} \\psi)^2 \\; , \\label{H} \\ee where $\\psi$ and $\\psi^{\\dagger}$ are canonically conjugate (field) variables, $m$ is the particle mass, and $g_4$ is the coupling that encodes the value of the scattering length.\\footnote{ In the context of dark matter, using (\\ref{H}) implies neglect of gravitational interactions, which in general is only possible up to some spatial scale. We discuss this question further in the conclusion.} In a macroscopic volume, a negative $g_4$ (i.e. an attractive interaction) leads to clumping, the process that had been investigated numerically in ref. \\cite{dew}. In the present paper we consider the case of a positive $g_4$ (a repulsive interaction). In the concluding section, we will also discuss briefly the case when the coupling is switched from positive to negative, as in a recent experiment \\cite{collapse}. We assume that the volume is macroscopic, and the gas is on average uniform.  A fundamental role in the kinetics of a Bose gas is played by the length scale corresponding to the wavenumber $k_{\\rm Bo}$, \\be \\hbar^2 k_{\\rm Bo}^2 = 4 g_4 m \\nave \\; , \\label{kBo} \\ee which we will call the Bogoliubov wavenumber. Here $\\nave = \\langle \\psi^{\\dagger} \\psi \\rangle$ is the average density of the gas. Particles with momenta $p\\gg \\hbar k_{\\rm Bo}$ can be considered as weakly interacting, while those with smaller momenta in general cannot. For $p \\ll \\hbar k_{\\rm Bo}$, elementary excitations of a fully condensed gas have a linear dispersion law; these are Bogoliubov's phonons \\cite{Bogoliubov}.  In addition to the Bogoliubov length $\\lambda_{\\rm Bo} = 2\\pi/ k_{\\rm Bo}$, which is time-independent, let us consider the time-dependent correlation length of the gas $\\lambda_0(t)$. The correlation length is defined as the characteristic scale at which the two-point function $\\langle \\psi^{\\dagger} (\\x) \\psi(0) \\rangle$ changes its behavior (e.g. from a power law to an exponential). The corresponding wavenumber is $k_0(t) = 2\\pi/ \\lambda_0(t)$. BEC corresponds to the growth of the correlation length with time, and we now summarize some previous results concerning this growth \\cite{Kraichnan}--\\cite{yale}.  If the initial value of the correlation length is small (we say that the initial sate is disordered), \\be \\lambda_0(0) \\ll \\lambda_{\\rm Bo} \\;, \\label{dis} \\ee BEC proceeds in two stages. At the first stage, $\\lambda_0(t) \\ll \\lambda_{\\rm Bo}$. If one assumes that the main role is played by interactions of particles with wavenumbers larger than or of order $k_0(t)$, one can use a quantum Boltzmann equation, because at this stage all such particles are weakly interacting. The solution to the Boltzmann equation \\cite{ST} shows that the main feature of the spectrum is a turbulent front that propagates to smaller wavenumbers. The position of the front can be identified with $k_0(t)$.  When the correlation length exceeds $\\lambda_{\\rm Bo}$, the particles begin to interact strongly with their local environment, and the standard Boltzmann equation cannot be used. At this second stage the growth of the correlation length continues through {\\em coarsening}, i.e. a process of alignment of the field on progressively larger spatial scales \\cite{KS,yale}, characterized by a power-law dependence of $\\lambda_0$ on time (at large times): \\be \\lambda_0(t) = {\\rm const} \\times t^{\\alpha} \\; . \\label{coar} \\ee We will return to discussion of the scaling exponent $\\alpha$ shortly. Also, as we will see below, the particle number contained in modes with $k<\\kBo$ is in fact concentrated in the ever-shrinking region $k\\sim k_0(t)$. So, at this stage the correlation length $\\lambda_0$ gives the typical size of correlated patches (i.e. regions over which the phase of the field is approximately constant).  If the initial value of the correlation length is already large (the initial state is partially ordered), \\be \\lambda_0(0) \\gg \\lambda_{\\rm Bo} \\;, \\label{ord} \\ee there is no room for a ``preparatory'' stage that can be described by a Boltzmann equation for the particles: one has to use a different description from the beginning. Here we consider only the late-time coarsening stage, which is present in either of the two cases.  How does one go about determining the scaling exponent $\\alpha$? The fastest the correlation length can possibly grow is at the speed of sound (cf. ref. \\cite{Kraichnan}), which in our case is the speed of Bogoliubov's phonons \\be v_s = (g_4 \\nave / m)^{1/2} \\; . \\label{vs} \\ee If one replaces (\\ref{H}) with a theory of non-interacting phonons, as it is done in the ``phase-only'' model of ref. \\cite{yale}, the only scaling law one can possibly discover is $\\lambda_0(t) \\sim v_s  t$, corresponding to $\\alpha = 1$. Although this value of $\\alpha$ is consistent with the results of numerical simulations \\cite{yale} of the full theory (\\ref{H}), the questions remain if the approximation of non-interacting phonons is a good one, and why their free propagation should have anything to do with the growth of the correlation length.  In the present paper, we attempt to answer these questions by taking a look at what happens in the momentum space, i.e. at the evolution of the power spectrum of the field. We find that a consistent description of the power spectrum is obtained by assuming that the longest-wave phonons, far from being non-interacting, are in the opposite, maximally nonlinear, regime. This assumption leads to a prediction of an inverse (i.e. directed towards the infrared) turbulent cascade of the particle number, with a definite value of the scaling exponent. This prediction is confirmed by numerical simulations in two and three spatial dimensions.  We stress that this {\\em strong turbulence}, characterized by strong quasiparticle interactions, is completely different from {\\em weak turbulence} found in a variety of physical systems on the basis of a Boltzmann equations for weakly interacting quasiparticles \\cite{ZLF}. Likewise, it is different from the turbulent solution obtained numerically in ref. \\cite{ST} for the Boltzmann stage of the evolution of a Bose gas.  Because phonons at wavenumbers $k \\sim 1/\\lambda_0$ are maximally nonlinear, the rate at which the particle number is transferred to still larger scales is of order of the phonon frequency $\\omega=v_s k$. The time at which the order is established at a scale $k$ is therefore $t \\sim 1/\\omega$, which leads to \\be \\lambda_0(t) \\sim v_s t \\; . \\label{lin} \\ee Thus, the linear scaling law for the correlation length obtains, and the speed of BEC is of order $v_s$, even though phonons are by no means free-propagating.  The paper is organized as follows. In Sect. 2 we present analytical considerations, and in Sect. 3 the numerical results. Concluding Sect. 4 contains estimates of condensation times for trapped atomic gases and galactic dark matter. ",
        "conclusions": "We now apply our results to two cases: (i) BEC in atomic vapors and (ii) formation of a Bose-Einstein condensate from galactic dark matter. In both cases, we estimate the speed of BEC as being of order of the sound velocity (\\ref{vs}). We can then estimate the time scale required for BEC to complete on physically interesting spatial scales.  Using the estimate $4\\times 10^{14}$ cm$^{-3}$, obtained in ref. \\cite{exp3} for the number density of sodium atoms, as $\\nave$ in eq. (\\ref{vs}) and noting that our parameter $g_4$ is identical to ${\\tilde U}$ used in that paper, we find $g_4 \\nave = 520 ~{\\rm nK}$, and \\be v_s = 1.4 ~{\\rm cm/s} \\; . \\label{vsgas} \\ee Thus, for a gas cloud a few micron across, we estimate that the coarsening time is shorter than 1 ms. The duration of evaporative cooling in ref. \\cite{exp3} was 7 s, so the coarsening time was not a limiting factor in that experiment. In the experiment with rubidium \\cite{exp1}, density was about two orders of magnitude smaller, which reduces $v_s$ by about an order of magnitude compared to (\\ref{vsgas}) ($g_4$ and $m$ are also different, of course, but by smaller factors). Still, the coarsening time remains much smaller than the time scale of evaporative cooling.  On the other hand, the relaxation time corresponding to the first, Boltzmann, stage of BEC in these gases is quite large. For a gas of lithium-7 atoms (which have an attractive interaction at low densities), with the parameters of the experiment of ref. \\cite{exp2}, this time was estimated in ref. \\cite{dew} to be of order 1 s. Similar estimates can be obtained for gases with repulsive interactions. First, form the dimensionless parameter of nonlinearity \\be \\xi_T = \\frac{4\\pi\\hbar^2 a n}{m T} \\; , \\label{xiT} \\ee where $a$ is the scattering length, and $T$ is the temperature in energy units. For rubidium-87, the authors of ref. \\cite{exp1} quote $a \\sim 10^{-6}$ cm, $n=2.5\\times 10^{12}$ cm$^{-3}$, and $T = 170$ nK. With these values of the parameters, $\\xi_T \\sim 10^{-2}$. In the Boltzmann approximation, the relaxation rate is proportional to $a^2$. Hence, we can roughly estimate the relaxation time as \\be t_r \\sim \\frac{\\hbar}{T \\xi_T^2} \\sim 0.4~{\\rm s} \\; . \\label{tr} \\ee The authors of \\cite{exp1} noted that after reaching the condensation temperature the condensate fraction grew for 2 seconds. Given how crude the estimate (\\ref{tr}) is, it seems plausible that this delay can be attributed to the Boltzmann relaxation.  The distinction between disordered and partially ordered initial states is important also for a Bose gas with an attractive two-particle interaction (a negative $a$). The wavenumber $\\kBo$ is now defined (in units with $\\hbar =1$) by \\be k_{\\rm Bo}^2 = 4 |g_4| m \\nave = 16 \\pi |a| \\nave \\label{kBo_attr} \\ee (the last equality is specific for three spatial dimensions). The quadratic dependence of the rate of collapse (fragmentation into individual drops) on the scattering length \\cite{dew} applies only when the initial state is disordered, i.e. the correlation length $\\lambda_0$ is smaller than $\\lambda_{\\rm Bo}$. In the opposite limit of a large correlation length, collapse in a sufficiently dense system will occur by linear instability. Indeed, the dispersion law for excitations with $k \\gg k_0$ is now \\be \\omega^2(k) = \\frac{1}{4m^2} \\left( k^4 - \\kBo^2 k^2 \\right) \\; , \\label{disp_attr} \\ee which corresponds to an instability for all $k < \\kBo$. The instability is the strongest for $k = \\kBo/ \\sqrt{2}$, where its rate is \\be |\\omega|_{\\max} = \\kBo^2 / 4m \\; . \\label{ommax} \\ee A large initial correlation length can be achieved, for example, if the scattering length was originally positive, but after some time was switched to a negative value, as in a recent experiment \\cite{collapse}. If we imagine that the gas is confined in a sphere of radius $R$ with Dirichlet boundary conditions for $\\psi$, and we gradually increase $\\kBo$, by increasing $|a|$ or $\\nave$, the instability will first appear when $\\kBo = \\pi / R$. Then, after $\\kBo$ passes a transitional region with $\\kBo \\sim \\pi / R$, the time scale of collapse will, as seen from (\\ref{ommax}), become inversely proportional to the scattering length and to the density of atoms. In our estimation, these scaling laws---especially the former---are well born out in the experimental results of ref. \\cite{collapse}.  For Bose gases considered in cosmological and astrophysical applications, one should take into account gravitational effects. If the particles are quanta of a real scalar field with a potential \\be V(\\phi)= \\half m^2 \\phi^2 + {\\lphi \\over 4} \\phi^4 \\; , \\ee then, in the absence of gravity, (\\ref{H}) is the correct nonrelativistic limit, provided \\be g_4 = \\frac{3\\lphi}{4 m^2} \\; . \\ee This limit is obtained by writing \\be \\phi = \\frac{1}{\\sqrt{2m}} \\left( \\psi \\e^{-imt} + \\psi^{\\dagger} \\e^{imt} \\right) \\ee and assuming that $\\psi$ changes with a frequency scale much smaller than $m$. Eq. (\\ref{vs}) for the sound velocity becomes \\be v_s = \\left( \\frac{3\\lphi \\rho}{4m^4} \\right)^{1/2} \\; , \\ee where $\\rho = m\\nave$ is the mass density.  Gravitational interaction (in Newtonian gravity) modifies the quasiparticle dispersion law, so that now \\be \\omega^2(k) = \\frac{1}{4m^2} \\left( k^4 + \\kBo^2 k^2 - k_J^4 \\right) \\; , \\label{disp_grav} \\ee where $k_J^4 = 16 \\pi G m^2 \\nave$ is the Jeans scale of the gas. (For a discussion of various length scales in a self-gravitating Bose gas, and various types of gravitational instabilities, see ref. \\cite{grinst}.) In a partially ordered Bose gas, eq. (\\ref{disp_grav}), similarly to eq. (\\ref{disp}), applies only at scales smaller than the correlation length $\\lambda_0(t)$. It follows that once $\\lambda_0$ exceeds a critical value $\\lmax$, \\be \\lmax^2 = \\frac{4\\pi^2 \\kBo^2}{k_J^4} = \\frac{3\\pi \\lphi}{4G m^4} \\; , \\label{lmax} \\ee the gas becomes gravitationally unstable via a linear (Jeans) instability and fragments into smaller clumps. Similar expressions were obtained in \\cite{it86}, \\cite{Peebles}--\\cite{Riotto&Tkachev} for the radius of a stable gravitationally bound object.  According to the general theory presented above, BEC at scale $\\lmax$ will occur through coarsening only if $\\lmax$ is larger than the Bogoliubov wavelength $\\lambda_{\\rm Bo}$, or equivalently $\\kBo > k_J$. This corresponds to \\be \\lphi > \\frac{4 m^3}{3} \\left( \\frac{\\pi G}{\\rho} \\right)^{1/2} = 10^{-25} \\left(\\frac{m}{1~{\\rm eV}} \\right)^3 \\; . \\label{bound_coars} \\ee In the last estimate, we have used $\\rho = 0.02 M_{\\odot} / {\\rm pc}^3$ \\cite{Firmani&al} for the mass density at a galaxy core. The time scale $\\tau$ required for coarsening to extend to scale $\\lmax$ can be estimated as \\be \\tau = \\frac{\\lmax}{v_s} = \\left(\\frac{\\pi}{G\\rho}\\right)^{1/2} =  2 \\times 10^8 ~{\\rm yrs} \\; . \\label{tau_coars} \\ee Thus, regardless of the coupling, $\\tau$ is much smaller than the current age of the universe (although much larger that the age at the time of galaxy formation).  On the other hand, the condition for the first, Boltzmann, stage of BEC to complete, obtained in \\cite{it91,Riotto&Tkachev} is, for the same value of $\\rho$, \\be \\lphi > 10^{-15} \\left(\\frac{m}{1~{\\rm eV}} \\right)^{7/2} \\; . \\label{bolt} \\ee If this condition is satisfied, completion of BEC on scale $\\lmax$ is guaranteed by (\\ref{tau_coars}). So, for the gas to be at least linearly stable on the scale of the galaxy core, $\\lmax$ should exceed the core diameter $2 r_c$. In view of (\\ref{lmax}), this requires \\be \\lphi > 3 \\times 10^{-4} \\left(\\frac{r_c}{1~{\\rm kpc}}\\right)^2 \\left(\\frac{m}{1~{\\rm eV}} \\right)^4 \\; . \\label{lam} \\ee Except for very small masses, (\\ref{lam}) is a stronger condition than (\\ref{bolt}) for all scales of relevance to galactic dark matter, $r_c > 1~{\\rm kpc}$. If $\\lphi$ satisfies (\\ref{bolt}) but not (\\ref{lam}), then the coarsening eventually sets off a Jeans instability, and the core fragments into smaller objects (Bose stars).  In conclusion, we have presented analytical arguments in favor of strong turbulence of the particle number occurring at the final, coarsening stage of Bose-Einstein condensation. We have found the form of the spectrum in the region of wavenumbers corresponding to an inverse cascade of the particle number. We have confirmed both the existence and the scaling exponent of the power-law spectrum by numerical simulations of classical fields in two and three dimensions. According to our theory, quasiparticles (Bogoliubov's phonons) are in the maximally nonlinear regime, implying that the correlation length grows linearly in time, with the speed of order of the speed of phonons. Finally, we have applied these findings to estimate the time scale of BEC in atomic vapors and in galactic dark matter.  I would like to thank I. Tkachev for getting me interested in theory of BEC and many useful discussions. Part of this work was done during the 1999 miniprogram ``Nonequilibrium Quantum Fields'' at the Institute for Theoretical Physics, Santa Barbara. I thank ITP for hospitality. This work was supported in part by the U.S. Department of Energy through Grant DE-FG02-91ER40681 (Task B)."
    }
}